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AD-A130 687 PROCEEDINGS OF THE STIP (STUDY OF TRAVELLING1/LAB HANSCOM AFB MA M A SHEA ET AL. 27 DEC 82
UNCLASSIFIED AFGL-TR-82-0390 F/G 3/2 NL
i/lllllI/llllEEEEEEEEmliElEhmhmhhhhhhhloEohhhhEEEohEEEEEEmhhIEmhhIEEEEmhhhEEohhohEEEEllEEEllllEEEEEEE~lEllEEEI
111W-. 1*28 12.5
1.4
MICROCOPY RESOLUTION4 TEST CHARTNATIONAL BUREAU OF STAtNOARS - ISS- A
ADA 03o68TAFGL-TR-82-0390SPECIAL REPORTS, NO. 233
Proceedings of the STIP Symposium on Solar RadioAstronomy, Interplanetary Scintillations and
Coordination With Spacecraft
Editors:
M.A. SHEAD.F. SMARTD.J. McLEANG.J. NELSON
27 December 1982
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This technical report has been reviewed andis approved for publication.
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SJUL 2 8 1983
COMPONENT PART NOTICE J 289
THIS PAPER IS A COMPONENT PART OF THE FOLLOWING COMPILATION REPORT: A(TITLE): Proceedings of the STIP (Study of Travelling Interplanetary Phenomena)
Symposium on Solar Radio Astronomy, Interplanetary Scintillations and
Coordination with Spacecraft Held at Narrabri, N.S.W., Australia on 28-30
November 1979.
(SOURCE): Air Force Geophysics Lab., Hanscom AFB, MA.
To ORDER THE COMPLETE COMPILATION REPORT USE AD-A130 687
THE COMPONENT PART IS PROVIDED HERE TO ALLOW USERS ACCESS TO INDIVIDUALLYAUTHORED SECTIONS OF PROCEEDINGS, ANNALS, SYMPOSIA, ETC- HOWEVER, THECOMPONENT SHOULD BE CONSIDERED WITHIN THE CONTEXT OF THE OVERALL COMPILATIONREPORT AND NOT AS A STAND'ALONE TECHNICAL REPORT.
THE FOLLOWING COMPONENT PART NUMBERS COMPRISE THE COMPILATION REPORT:
AN: Pool 435 TITLE: Radio and X-Ray Evidence for Electron Acceleration inSolar Flares.
PO01 436 Magnetohydrodynamic (MID) Description of Solar RadioBursts.
PO01 437 Profiles of Coronal MID (Magnetohydrodynamic) DisturbanceAssociated with the Solar Flare on April 27, 1979.
PO01 438 Fine Structure in Metre-Wave Solar Radio Bursts.PO01 439 The Mechanism of the U-Shaped Spectrum of Type IV Solar
Radio Bursts.POOl 440 The Proton Flare of 1976 April 30.POOl 441 Relativistic Solar Particle Events during STIP (Study
of Travelling Interplanetary Phenomena) Intervals II
and IV.POOl 442 Thermal Cyclotron Radiation from Solar Active Regions.POOl 443 Observations of the Quiet Sun at Decametre and Metre
Wavelengths.PO01 444 The Rise, Decline, and Possible Relationship between
two Flare Producing Regions Including Details of theActivity Patterns of Their Known Antecedent and DescendantPlages.
POOl 445 Theoretical MHD (Magnetohydrodynamic) Simulations ofCoronal Transients and Interplanetary Observations.
PO01 446 Three-Dimensional Distribution of the Solar Wind VelocityDeduced from IPS (Interplanetary Scintillation)Observations.
POOl 447 Interplanetary Scintillations at 103 MHz.POOl 448 Analysis of IPS (Interplanetary Scintillation) Data Using
a Periodic Phase Screen Model for the Solar Wind.POOl 449 STIP (Study of Travelling Interplanetary Phenomena) VII
' I Observations of the Northern Solar Corona - A First ReportA ' on the Scope of Data and the Early Results.
POOl 450 Unusually Intense Jovian Decametric Emission Observed on
it t rvi1979 March 7, 1920 -Til 200tr. document has beon c.-,p--vcdiUr plblic release an6ii!a 'at its
~ f t I i-A; tri . -u-u i.- unlimited.
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AFGL-TR-82-0390 j.' ~ ,g'4 TITLE (ond &,ibiliI.) S. TYPE OF REPORT & PERIOD COVERED
PROCEEDINGS OF THE STIP SYMPOSIUM ONi SOLAR RADIO Scientific, Interim.ASTRONOMY, INTERPLANETARY SCINTILLATIONS AND _____________
COORDINATION WITH SPACECRAFT B. PERFORMING ORG. REPT NUMBERSR, No. 2337
7.~ S* . CONTRACT OR GRANT NUMBER(aJ
Editors: M. A. Shea, D. F, Smart, D. J, McLeanand G. J. Nelsonw*
9. PERFORMING ORGANIZATION NANE AND ADDRESS IS. PROGRAM ELEMENT. PROJECT. TASKAREA A WORK UNIT NUMBER"S
Air Force Geophysics Laboratory (PRIG) 61 102FHanscom AFB, Massachusetts 01731 23l1IO8
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Air Force Geophysics Laboratory (PUG) 27 December 1982
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ISl. SUPPLEMENTARY NOTES
Division of Radiophysics, CSIRO, P. 0. Box 76, Epping, N.S.W., 2121, Australi**CSIRO Solar Radio Observatory, Culgoors, P. 0.* Box 94, Narrabri, N.S.W. 2390,
Australia
Solar physic s Solar particles
Solar radio astronomy Interplanetary mediumInterplanetary scintillations Cosmic raysSolar wind Jupiter radio emission
20. ABSTRACT (Continue o teoo side.... finocesry, end Identify by block rtteeb..)
This report is a compilation of 16 papers presented at the Study ofTravelling Interplanetary Phenomena (STIP) Symposium on Solar RadioAstronomy, Interplanetary Scintillations, and Coordination with Spacecraft,held 28-30 November 1979 at the Culgoora Radio Observatory in Narrabri,N.S.W., Australia. These papers span a range of multi-disciplinary studiesof various solar emissions propagating from the sun through the interplanetamedium to the earth and beyond to Jupiter. This includes papers on'solarradio emissions. solar region evolution; solar particle propagation; the
DD I JAN73 1473 EDITION OF I NOV 65 IS OBSOLETE Ucasfe
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modeling of solar induced disturbances in the interplanetary mediumand theuse of radio measurements to probe the turbulence of the interplanetarymedium. Theoretical studies and analyses, computer simulation and in situmeasurements of these phenomena are included. The subjects covered aresolar physics, solar radio astronomy, interplanetary scintillation measurement,cometary studies, direct spacecraft measurements and energetic solar particlepropagation in the interplanetary medium.
Unclassified
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Dist sp"-cial
t
Foreword
The Study of Travelling Interplanetary Phenomena (STIP) is one of the inter-disciplinary studies established by the International Council of Scientific Union'sScientific Committee on So re Slar rstral Physics (SCOSTEP). The scientific ob-jectives of STIP are the study of phenomena that propagate through the interplane-tary medium and a search for understanding of these phenomena during both quiet(i.e. normal) and active periods.
About 300 scientists are actively participating in STIP, their interestsranging from solar physics (insofar as it concerns the initiation of phenomenawhich noves out from the Sun) to the observation and study of comets and planetarymagnetospheres and ionospheres. Solar radio astronomy and interplanetary scintil-lations of discrete radio sources, solar and galactic cosmic rays, solar windplasma and interplanetary magnetic fields, and interstellar interactions are amongthe topics of interest.
The STIP Symposium on Solar Radio Astronomy, Interplanetary Scintillations and
Coordination with Spacecraft was held in November 1979 at the site of the CSIROCulgoora radioheliograph in Narrabri, N.S.W., Australia. This specialized sym-posium, co-sponsored by the IAU, COSPAR, and IAGA, was attended by -35 scientistsfrom 10 countries. A total of 21 papers were presented allowing ample time fordiscussion. In addition, a tour of the radioheliograph and other on site CSIROradio astronomy equipment was conducted by the CSIRO personnel.
Sixteen of the papers presented at the symposium are contained in this proce-edings. Although some of the manuscripts were prepared in camera-ready format bythe authors, most of the papers were edited and re-typed after review by the refer-rees. The editors are most grateful to Ms. Marie Vickery for her careful reviewand typing of eight of these manuscripts. The editors also thank Ms. KarenFlaherty, Ms. Anne A. Bathurst, Ms. Louise C. Gentile, and lis. Stacey A. Horningfor their assistance in typing and reference searches.
Finally, on behalf of all the participants, we wish to thank the CSIRO groupfor their work in providing the conference facility and the local arrangements for
this symposium.
M. A. SheaD. F. SmartD. J. McLeanG. J. Nelson
iii
... .... .. .. . . . . .. .
Contents
Foreword iii
1. Radio and X-ray Evidence for Electron Acceleration in Solar FlaresR.T. Stewart
2. tiagnetohydrodynamic (M'HD) Description of Solar Radio Bursts 19S.T. Wu
3. Profiles of Coronal MHD Disturbance Associated with the Solar Flare on 37April 27, 1979
R.V. Bhonsle, S.S. Degaonkar, and S.K. Alurkar
4. Fine Structure in Metre-Wave Solar Radio Bursts 43D. rcConnell and G.R.A. Ellis
5. The Mlechanism of the U-shaped Spectrum of Type IV Solar Radio Bursts 49Li Chun-sheng, Zheng 11ing-wu, and Yao Jin-xing
6. The Proton Flare of 1976 April 30 65
S.I. Avdjushin, N.K. Perejaslova, V.V. Fomichev, and I.M. Chertok
7. Relativistic Solar Particle Events During STIP Intervals II and IV 71t.A. Shea and D. F. Smart
8. Thermal Cyclotron Radiation from Solar Active Regions 85V.V. Zheleznyakov and E. Ya. Zlotnik
9. Observations of the Quiet Sun at Decinetre and Mfetre Wavelengths 97A.A. Abranin, L.L. Bazeljan, and R.A. Duncan
10. The Rise, Decline and Possible Relationship Between Two Proton Flare 107Producing Regions Including Details of the Activity Patterns oftheir Known Antecedent and Descendant Plages
Susan M.P. McKenna-Lawlor
v
Contents
11. Theoretical MHD Simulations of Coronal Transients and Interplanetary 137Observations
M. Dryer, R.S. Steinolfson, and Z.K. Smith
12. Three-dimensional Distribution of the Solar Wind Velocity Deduced 153
from IPS ObservationsT. Kakinuma, HI. Washimi, and Me Kojima
13. Interplanetary Scintillations at 103 MIlz 163S.K. Alurkar and R.V. Bhonsle
14. Analysis of IPS Data Using a Periodic Phase Screen Model for the 167
Solar WindA. Basu and S.K. Alurkar
15. STIF VII Observations of the Northern Solar Corona--A First Report on 183the Scope of Data and the Early Results
Thomas A. Croft
16. Unusually Intense Jovian Decametric Emission Observed on 1979 March 7, 2091920-2040 UTK. Bullough and W. Gibbons
vi
THE CULGOORA RADIOHELIOGRAPH
This unique instrument consistes of 96 antenna (each antenna being ateerable paraboloid about 45 feet in diameter) arranged in a circular
array 1.86 miles in diameter. This radioheiograph has been the worldstandard for meter wave length solar observations since it began oper-ation in 1967.
Vrii
PROCEEDINGS OF THE STIPSYMPOSIUM ON SOLAR
RADIO ASTRONOMY,INTERPLANETARY
SCINTILLATIONS ANDCOORDINATION WITH
SPACECRAFT
1. Radio and X-Ray Evidence forElectron Acceleration in Solar
Flares
R. T. StewartDivision of Radiophysics, CSIRO,
P.O. Box 76, Epping, N.S.R. 2121, Australia
Abstract
Radio and X-ray observations, as well as 4-ray and interplanetary electron ob-servations, are briefly reviewed for evidence of elect,in acceleration during solarflares. The evidence collected so farsuggests that in some large flares a secondstage of particle acceleration takes place where ions and electrons are acceleratedto very high energies. This second stage follows no more than several minutes afterthe first or impulsive stage, which occurs during the initial minute or so of thesolar flare. During the impulsive phase electrons are accelerated to energies-10-1O0 keV, probably by bulk energization, where stored magnetic energy is convertedInto thermal energy. Energy-loss arguments favoar the electrons having a quasi-thermal (near-Maxwellian) electron velocity distribution. During large flaressecond-phase electrons are accelerated to much higher energies, -i0 MeV, and appearto have a power-law velocity distribution. The second-stage acceleration mechanismmay be associated with a coronal shock wave, but so far there is no clear evidencethat electrons with energies "1100 keV are directly associated with a shock wave.
I. INTRODUCTION
The problem of particle acceleration in solar flares has been studied for manyyears now but is still not fully understood. Evidence for electron accelerationcomes from direct measurements of streams of high-energy electrons in interplanetaryspace and from indirect manifestations, such as burst emission from the Sun atradio, EUV, X-ray and y-ray wavelengths. Radio bursts at metre wavelengths appearto be due mainly to plasma radiation excited by electrons with energies -10-100 keVand those at microwave wavelengths to gyro-synchrotron radiation from electronswith energies l100 keV to 1 MeV. Hard X-ray bursts appear to be due to bremsstrah-lung from electrons with energies -10-100 keV and y-ray continuum to bremastrahlungfrom relativistic electrons with energies -1-10 MeV.
Because energetic ions -10 GeV and relativistic electrons 10 MeV are ejectedfrom the Sun during large flares it is generally believed that these high-energyparticles are accelerated in a second-stage process. During the first stage, whichoccurs during the impulsive phase of the flare, electrons are accelerated toenergies -10-100 keV. Most of the accelerated electrons have energies 30 keV andit has been suggested that these first-stage electrons might come from a quasi-thermal source region where the electron temperature is -2-4 x 108 K. In thesecond stage electrons and ions are accelerated to relativistic energies and appearto have non-thermal power-law distributions. It is thought that some form ofstochastic process such as Fermi acceleration is responsible.
2. FIRST-STAGE ACCELERATION
This occurs during the first 10 s to I min of a solar flare. The radio signa-ture is a group of Type III bursts at metre and decimetre wavelengths and animpulsive microwave burst. Hard and soft X-ray bursts, Ha and EUV flashes alsooccur. Sometimes fine structures with durations -I s are observed in the microwaveand hard X-ray bursts and seem to be associated with individual Type iI bursts atmetre and decimetre wavelengths (Figure 1). This suggests that the accelerationprocess is very short (I s) and repeated. Impulsive bursts are also closelyassociated with interplanetary electron streams with energies -10-100 keV but notwith proton streams (Lin, 1974). Only a small fraction of the electrons acceleratedin the lower corona near the hard X-ray source region escape to interplanetary space(Lin and Hudson, 1971), giving rise to Type III and Type V radio bursts as theytravel through the corona (Figure 1); the remaining electrons are either trappedon closed magnetic field lines, where they give rise to microwave bursts by gyro-synchrotron emission, or are lost by collisions and give rise to hard X-ray burstsby bremsstrahlung.
Most of the flare-accelerated electrons seem to be associated with emission ofthe hard X-ray burst. For example, Hoyng et al. (1976) estimate that -1036electrons per second with energies >25 keV are needed for a non-thermal electrondistribution with spectral index -4 to explain a strong impulsive hard X-ray burstby thick-target bremsstrahlung. The corresponding energy input for a burst lasting10 s is -I03 ergs, which is an appreciable fraction of the total flare energy.The reason for the high-energy requirement in non-thermal electrons is that most oftheir energy is lost by collisions in the target area and only a small fraction-10
-5 goes into bremsstrahlung (Hoyng et al., 1976). This had led a number of in-
vestigators to consider thermal or quasi-thermal source models for impulsive hard
2
25
50
100 -
600go ARCH 21, '1972
700
G o-
S30
2 24-34keV -40
20 -- 0
10 -
10 0
0112 0113 0114
- TIME (UT)
Figure 1. Example of a metre-wavelength Type Ill-V burst(upper) associated with a hard X-ray burst (lower). Thelower figure is a copy of Figure 2 of Hoyng (1975). TheType III-V burst was recorded by the Culgoora radio spectro-graph with a time resolution of 0.25 s. The hard X-rayburst was recorded by the Utrecht hard X-ray spectrometer onboard the ESRO TD-IA satellite with a time resolution of1.2 s. The energy ranges of the five X-ray channels and theobserved frequency range of the radio burst are indicated(Stewart, 1978).
3
X-ray bursts, In these models the electrons are impulsively heated in a smallsource region and relax to a near-Maxwellian distribution where the electron
temperature Te is -1- x LO8 K and the ion temperature Ti is considerably less
(Te z 10 T) (Brown et al., 1979; Ramaty et al., 1979; Smith and Lilliequist,
1979). Considerable energy is lost from the source region by conduction to thesurrounding plasma but even so the models are considered to be 10-100 times moreefficient in producing X-rays than the non-thermal thicK-target model.
Some observational support for the thermal model has been presented by
Crannell et al. (1978), who show that the impulsive hard X-ray burst profile oftenhas a symmetrical rise and fall (Figure 2) which could be caused by a reversible
0
40f
20
4 0 00UNIVERSAL rIME
Figure 2. The time profiles of the sum of theX-ray flux in channels 2 (28-55 keV) and 3
(55-82 keV) of the OSO-X-ray spectrometer (NASAGoddard Space Flight Center) and the 10.5 GHzmicrowave flux (Institute of Applied Physics,
University of Bern). The X-ray profile has beennormalized to the same peak height as the micro-wave profile (Crannell et al., 1978).
process such as adiabatic compression (heating) and expansion (cooling) of a smallsource region (Matzler et al., 1978). They also show that impulsive microwavebursts are closely correlated with hard X-ray bursts (Figure 2). This suggests thathard X-rays and microwaves are emitted from the same source of thermal electrons.Combining the two sets of observations, Matzler et al. (1978) and Dulk et al. (1979)find electron temperatures Te 108 K, emission measures n
2V _
104
5 cm
-3 , source
sizes A k1018 cm, electron densities ne 1 Q19 cm- 3
, and magnetic fields with
B -
100-400 G (Figures 3a and 3b). We note that the electron and ion densities
derived from these observations are considerably less than the values -l01 cm-3
used in the theoretical models mentioned above.
4
-...... T-EM I
10 . ,*
T8 =1.5,EM 4 9"0S
Te s3'0,EM45' 1.0
I(E)
'\
0.1 %S
001
1 * I * I . I
20 40 60 80100 200
E (kev)
Figure 3(a). X-ray spectrum of the impulsive burstof 1972 May 18, 14ho6m09s UT. The datum points arefrom Hoyng and Stevens (1973). The solid and dashedcurves represent X-ray spectra derived from electronthermal distributions with the parameters indicated onthe figure. T8 is the electron temperature in 108 Kand EM4 5 the emission measure in 104 5 cm- 3 (Dulk etal., 1979).
So far there have not been routine spatial observations of hard X-ray bursts.It is hoped that during the solar maximum mission such observations willsettle the question of whether the X-ray electrons are thermal or non-thermal andwhether these electrons are located primarily in the chromospheric-corona transitionregion, where the ion density ni is -1011 cm
- 3 , or in the low corona, where ni is-1(9 cm-3 . A recent stereo-spacecraft observation reported by Kane et al. (1979)indicates that the coronal contribution for one hard X-ray impulsive burst is 600times less than the chromospheric contribution.
Lin and Hudson (1971) and Stewart (1978) have estimated that 1033 to 1035
electrons with energies >10 keV escape from the flare region during the impulsivephase. If the total number of accelerated electrons is represented by the hardX-ray burst then as many as 1038 non-thermal electrons or >1036 thermal electronsare accelerated during the Impulsive phase. It would appear that <10% thermal or<1% non-thermal electrons escape to the higher corona to produce the Type ll-Vburst and the interplanetary electron event.
5
SI- . ,= ! . sI . I ' I
400
66 •
200
<2 SFU. • L I , I II I
100
80
60/40 L
20 J
S<2 SFU
2 4 6 810 20f (GHz)
Figure 3(b). Microwave spectrum of the event ofFigure 3(a). The datum points are from Hoyng andStevens (1973). The solid curve is a hand-drawn fitto the datum points. The dashed line was computed byDulk et al. (1979) using the parameters L8 and B2indicated on the figure, where L8 is the source sizein 108 cm and B2 the magnetic induction in 102 G.T8 - 2.3 and EM4 5 - 2.0 were derived from Figure 3(a)(Dulk et al., 1979).
The impulsive phase electrons are thought to be accelerated by some form ofbulk energization where stored magnetic energy is converted into electron energyduring the flare. Various processes involving the acceleration of electrons byinduced electric fields caused by magnetic field line motion have been reviewed bySmith (1974) and Ramaty et al. (1979).
3. SECOND-STAGE ACCELERATION
In very large flares there appears to be a second stage of acceleration whereelectrons reach energies of -10 MeV and protons and ions even higher energies,-10 GeV. Interplanetary particle observations suggest that energetic ions and
6
relativistic electrons are released from the Sun a few minutes after the 10-100 keVimpulsive-phase electrons are accelerated (Lin, 1970). Microwave and hard X-raybursts show a distinct first and second phase (Frost and Dennis, 1971) with theX-ray spectrum becoming harder in the second phase corresponding to an increase inthe proportion of high-energy electrons as second-phase acceleration takes place(see Figure 4 and also Hoyng et al. (1976)). Gamma-ray observations of the 1972
t, o N UCN56
0. .. ot.a.d
Figure 4(a). Time profiles of thehard X-ray flux for differentenergy ranges during the two-stage
event of 1969 March 30 (Frost andDennis, 1971). The second stage
began near the starting time of aType II burst indicated by the
vertical arrow.
August 4 flare (Chupp et at., 1975) indicate that the relativistic electrons(E - 1-10 MeV) producing *-ray continuum (0.35-8 MeV) by electron-electron andelevtron-ion bremsstrahlung and the high-energy protons (E - 10-100 MeV) producing
r-ray line emission (2.2 MeV) by nuclear reactions are accelerated several minutesafter the 10-100 keV impulsive phase electrons (Figure 5) (Bai and Ramaty, 1976).Both escaping and trapped flare electrons appear to have power-law spectra up toenergies -10 MeV (Lin et al., 1978; Bai and Ramaty, 1976).
The high proton-to-electron ratio (at, say, 10 MeV) observed in large flaresis consistent with the particles being accelerated in a stochastic process such as
7
0'. .T .... ....
to'I
" " '
S10<
o0 +
i
10 a20 40 60 100 200 JOU 20 40 60 100 200 300
ENERGY Kev,
Figure 4(b). The hard X-ray photon energyspectrum at two times during the impulsiveburst (a) and during the second-stage burst(b) of 1969 March 30. The impulsive burstspectrum shows a steepening above -100 keVwhich is fairly typical of flash phaseevents. The second-stage spectrum isharder than the first stage and fits asingle power law up to the maximum energyof the detector. (After Frost and Dennis(1971) and Lin (1974).)
Fermi acceleration (Ramaty et al., 1979). The energy gained by repeated reflectionsoff magnetic field irregularities is greater than the energy lost by collisions forelectrons with energies >1100 keV and for protons with energies 0.5 keV (Ramatyet al., 1979). Hence a high-energy tail is required for the inital electron distri-bution if Fermi acceleration is to occur. This high-energy tail could be providedby the first-stage acceleration of electrons during the impulsive phase ol theflare. Protons, on the other hand, could be accelerated from thermal energiesdirectly during the second stage (Ramaty eL al., 1979).
4. EXTENDED MICROWAVE BURSTS
Stewart and Nelson (1979) have found that the second-stage microwave burstis closely correlated with the hard X-ray burst (Figure 6a). These extendedmicrowave bursts have a fMatter spectruim (below the turnover frequency) thanimpulsive bursts (compare Figures ith and hb) and this can he ,xplained byassuming that the radiating e l ectrons are trapped in a non-unifiorm magneticsource region such as a dipole magne tic liild. III Sulch a source similar numbersof electr'ons with energies -1()0 keV are reqlired to cexplain the microwave and hard
8
4 AUGUST, 1972
2 9 -41 k4V(photons cm
2 sec1)
* (van Beek et al 1973)
I 0.35-8MeV
(photons cmfi2see')(Chupp et al
1975)-
UW Z3 GHzn (watts m'2 Hz'l) -
0_ ~(Croomn and""\ X Harris 1973)
S2.2 MeV
(photons cm 2 sec- ) ,
(Chupp et al 1975)_
4'i
6h2Off 6h30m 6h401"
TIME (U T)
Figure 5. Time dependences of radiations of theflare of 1972 August 4. The three upper lines arethe measured time profiles of X-rays (29-41 keV),gamma rays (0.35-8 MeV), and microwaves (37 GHz).The error bars in the lower part of the figurerepresent the measured intensities of the 2.2 MeVline. The solid, dashed, and dotted lines arecalculated time profiles of the 2.2 MeV line. The
solid line is obtained by assuming that the in-stantaneous number of nuclei in the flare region
has the same time dependence as that of theobserved 0.35 to 8 MeV gamma rays. The dashed anddotted lines are obtained by assuming that thetime dependence of the nuclei is the same as thatof the 29 to 41 keV X-rays. (For the solid and
dotted lines the photospheric 3He abundance ratio3
He/H equals 5 x 10-5; for the dashed line,3He/H equals 0.) Note that the solid line fits
the 2.2 MeV data better than the dashed or thedotted line. Hence both the 1-10 MeV electronsemitting the 0.35-8 MeV y-continuum and the10-100 MeV protons responsible for the 2.2 MeVy-line are accelerated several minutes after the10-100 keV impulsive hard X-ray burst electrons.(After Bai and Ramaty (1976).)
9
10"
*-375GHZ
9-4GHZZ0 7 d .. (7)
10P
75GHAEFUXD NZT(Z
2 Z
10,,
o 1& IC 0 (5) -(2
10;C 0
MICROWAVE FLLD< DENSITY(*" Wr62 HZ-)
Figure 6(a). Observed correlation between the
100 keV hard X-ray and 3.75 Gliz or 9.4 GHz
microwave flux densities at the time of X-ray
maximum in nine extended (second-stage) flare
events. The full line shows the least-mean-
squares fit (logarithmic) to the 9.4 Glz data
and the short dashed line the fit to the 3.75
GHz data. The nine events are listed in
Table I of Stewart and Nelson (1979).
10tA
5
\ '"A25 /
S(3)
U" 05 () -\ (2)
(c) - (3)
025-(
(6)-* (7)
.~ 2 \\ * (ez (...
025 05 1 2 5 10 20 40 60
FREQUENCY (GHz)
Figure 6(b). Microwave spectra at the time of
100 keV hard X-ray maximum for the nine eventsof Figure 6(a). Each spectrum is normalized to
the 3.75 Gliz value (Stewart and Nelson, 1979).
10
X-ray emission (Nelson and Stewart, 1979; Takakura, 1972). The turnover fre-quencies of extended microwave bursts are usually -i0-20 GHz (Figure 6b) and indicate
maximum values for the magnetic induction B e 500-1500 G (Nelson and Stewart, 1979).With the above source model we expect to find larger source sizes and strongermagnetic fields for second-stage microwave and hard X-ray bursts than for impulsivephase events. We also expect the size of the microwave burst to increase with de-creasing frequency. At present there are not enough high spatial resolutionobservations to test this model.
5. RADIO EVIDENCE FOR SHOCK WAVE ACCELERATION
Second-stage stochastic acceleration of particles in shock waves has been con-sidered by several authors, including Schatzman (1963) and Gubchenko and Zaitsev(1978). The clearest evidence for the presence of a shock wave in the corona duringa solar flare is the Type ii radio burst (Wild et al., 1953). These bursts usuallyoccur during proton flares (Van Hollebeke et al., 1975) near the start of thesecond-stage hard X-ray event (see arrow of Figure 4a). The Type 11 burst is thoughtto be due to plasma emission from mildly relativistic electrons (E 20 keV)accelerated by the shock wave. Sometimes these electrons escape from the Type 11shock wave and produce herringbone features, which are rib-like bursts with fast-frequency drift rats similar to those of Type III bursts. The high circularpolarization .f fundamental herringbone bursts (Suzuki et al., (1980)indicate thatthe electron streams travel along ordered magnetic field lines. Electron streamsof this kind have also been observed by spacecraft to travel along magnetic fieldlines away from the Earth's bow shock (R.P. Lin, 1979 - personal communication) andaway from interplanetary shock waves (Richter et al., 1979).
Type 11 shock waves probably also accelerate electrons to produce flare con-tinuum (FC II) bursts at metre wavelengths (Robinson and Smerd, 1975). The highbrightness temperature observed in some of these bursts, >101 K (Duncan et al.,1980; Magun et al., 1975), suggests that the flare continuum is coherent plasmaemission from electrons with energies l0 keV. According to Dulk et al. (1978), aflux of at least lO
3 electrons per second is required in the flare continuum source
region located at a height .6 R0 above the photospherc. For a typical burst of-l000 s duration this means that -lO
31 ' 0o L103 electrons with energies -0 kecV must
te accelerated by the shock wave as it passes through the corona, assuming that theelectrons are not trapped in the source region but escape over times l s. Thisnumber is considerably less than the value -1O1' required for the 10-100 keVelectrons in the extended hard X-ray source region.
It was originally thought that moving Type IV bursts were due to synchrotronemission from relativistic electrons with energies >1 MeV (Boischot and DInisse,1957). The possibility of large populations (_10
33 particles) of high-energy
electrons at the large source heights of moving Type IV bursts created a con-siderable interest in these events. However, later observations with the Culgooraradioheliograph showed that moving Type IV bursts at 80 and 160 ,MIz were oftenpolarized in the circular sense. This led Bulk (1973) and others to consider modelsin which the radio emission was attributed to low-harmonic gyro-synchrotron radia-tion from mildly relativistic electrons (-10-100 keV). More recent observations,especially at 80 and 43 MHz, have shown that the moving Type IV burst in its earlystages can have brightness temperatures -'1010 K (Figure 7) (Stewart et al., 1978a;Duncan et al., 1980). Such high brightness temperatures, together with tht obst rvedcircular polarization, suggest that the moving Type IV burst, at least in its earlystage of development, is due to coherent plasma emission from 10-100 ktV lectrons.
11
There is little direct evidence to suggest that these electrons are accelerated byshock waves. Indeed, the close association found between eruptive prominences andmoving Type IV sources (Stewart et al., 1978b) would seem to exclude the possibilityof shock wave acceleration.
91015L ib
5, S~ OL500L 9 0L
90LL
",--\L ' "t \ 5L'0
B \',,70L
40 80 160 4,0 80 160
-JK . R 1 d)- w
~30R
\ \".2 \ ' \
70R
7k
~ 70R
40 80 160 40 80 160
FREQUENCY (M,,Hz)}
Figure 7. Observed source brightness temperatureand degree of circular polarization of moving TypeIV bursts at 160, 80 and 43 MHz. Source heights(solar radii above the centre of the Sun) and times(UT) are shown in the following table.
12
- I I I mmmmmm .. _ - ....
Table to Figure 7
(1) -(iv) -(ii) (v)
(iii) (vi) .........
(vii) ...
(a) 1975 Aug. 22 (i) 0 1h30 m at 1.4 Re;
(ii) 0jh40m at 1.7 Re; (iii) 02ho0 m at 2.4 Re;
(iv) 0 2h20m at 2.7 Re; (v) 0 2h4Om
at 3.0 Re.
(b) 1977 Oct. 4/5 (i) 2 3h50m at 1.3 R3;
(ii) 0 0h0 8m at 1.7 Re; (iii) 00 h4 6m at 2.2 Re;
(iv) olh 1 3m at 2.5 Re; (v) olh 3 0 m at 3.0 Ro .
(c) 1977 Sept. 20 (i) 0 3h1 4m
at 1.3 Re;
(ii) 0 3h19m at 1.4 Re; (iii) 0 3h25m at 1.6 Re;
(iv) 0 3h35 m at 1.8 Re; (v) 0 3h40m at 2.1 Re;
(vii) 0 3h50 m at 2.6 Re; (viii) 0 4 ho0m
at 3.1 Re.
(d) 1977 Oct. 3/4 (1) 2 3h35m at 2.1 Re;
(ii) 2 3h51 m at 2.3 Re; (iii) 0 0 h0 3m at 2.6 Re;
(iv) 00 h24m at 3.0 Re.
L and R refer to LH and RH circular polarization.At 4) MHz the degree of circular polarization isaccurate to ±5% for 0-15% and ±10% for 15-50%. Themaximum source velocities are (a) 400, (b) 300,(c) 600 and (d) 400 km s
-1. Only the later part of
event (d) was recorded.
6. CONCLUSION
Radio and X-ray observations indicate that the bulk of the electrons accele-rated during the impulsive phase of the solar flare have energies in the range10-1000 keV. During very large flares there is evidence from interplanetaryparticle measurements and from y-ray observations that electrons are accelerated tomuch higher energies, IQ 0 GeV. It is still not clear how the particles areaccelerated during these two stages. A popular idea is that during the impulsivephase the electrons are impulsively heated to very high temperatures, -108 K, bysome form of bulk energization where stored magnetic energy is converted intothermal energy. These hot electrons then relax to form a quasi-thermal (near-Maxwellian) distribution and produce the hard X-ray burst near the base of thecorona by bremsstrahlung. A few per cent of the hot electrons then escape to higherregions of the corona to produce Type III-V bursts and interplanetary electronevents. Shock-wave acceleration is still an attractive possibility for the second-stage process, although there is no clear evidence to show that a coronal shock waveis capable of accelerating electrons and ions to the very high energies attained bysecond-stage particles.
IIII
References
Bai , T . , and Ramat y , R. :19/1u , "Gda -a-r ay andi mi c rowavc o-V I dL-11 V tI L w-' pulisOS of
acceleration in solar Ilares', Si'1 .rli\~5. ,9, 14 i.
Boischot, A., antd Ilenisse, J.: 19),, "'S rij' d,- tv,1% I t. 1 '0 1 iL2i III Jii-
rayons cosmiques aSS[iJli i S erupt ionsl iliriiiii'spiuiri~pi' -S S.Aad i I.Par is,
245, 2194.
Brown, J.C., Meirose-, fl.B.. and Spiev, 1).S.: 1979, ''IrOdoct il~j 11In it I ais it, n Ivs
conduction froint by rapid coronal heat ing aind its rou fin sol ar h-trd X-rac
bursts", Astnoiphy:;. J., .228i, ')92.
Chupp, h..., Forrost L, D).J., and Siri, A.N.: 1975, ''high energy glt-a'radiationabove 300 ke\' a.,,idt(-d siLI. Sol ar act ivit N', inl Sl I I, ( Ima and C'I
Radiations (Proc . lAU Symp. ',8) (ed . S . Kanie) , p). 341 lRce I , 'X',rdrioht
Crannell1, C.J. , Frost , K.J. , Utzlur, C. , Ohki, K. , and Saba, .1.:"Impulsive sol ar X-ray ibor~ts , Astrophvs. J., 223, 62(0.
Croom, 1). L..,and Hlarris, L.. I). :1971, ''Solar radio 'vc-nts a) 1,, -1 tIi; inthe period 20 julv - 14 August 1972'' in Colle cte d Da ta N , lo r t In )It .:c.197-2 Sotar-Ter restLr id a vents (ed leIe F. Colfey I ). I'A(;-?8. Part 1,p 'I0
(IS. l)epa rtmnn I f merce, NOAA , Bon ulerI
Dulk, G .A. : 197i3, ''The gyro-synchrot ron rad iat 1i)) I rim TI'''ifn,, tl % V -ur, -~ i
the solar coronai, Sol I Pjyj, 32, 491.
Dulk, G.A. , Mel rose , 1. e., and Sme rd, .F.: 197S, ''l4 111 ' i IlloIll 1,: .r t idistributio tll t ions'', Proc. Ast ron. So-: .Auis t. , ,.
lulk, G.A. , Mel rose, 1.1., antd Wlhito-, "!ht.97,'Li,5r..'i I.'.:lOi
from qoasi-therrsal electrons and~ appli~ations to solar ilar,-s'', A~D~.234, 1137.
Duncan, R.A., Stewart, K.t. , and Nels-n, 11.,;.: 1980, ''on. r\ ri ht 1%p, IV'solar mtre-WaVV radio e-missions', in Solar and otrlanIra-
IALD Symp. 91) (eds. M. Dryer and E. Tandberg-llanssen) p. 391 (Reldel, Dordrecht).
Frost, K. 3., and Dennis, B.K. : 1971, "ELvi dence from Duirl N-ra !, tparticle acceleration in -1 sl-ar flaro'', As.t~ 2 p . , 16)
Gubchenkn , V.M., alill Zaitsev, V .V.: 1978, 'on lprtL'n ind , It- twn ' I a nShOok waves during Ilarge solar flares"', l'roc . Ast Lran. -S'c . A\),"I
Iloyng, P.: 1975, ''StUdies in hard X-ray emission from solai [Iar,'s and) n1 ov, lot ro n
radiation front a cold lmanet )-lsaPh.D. Thiesis, Lniv, rjty of tti- tit.
Iloyng, P. , and Stevens, CG.A. : 1973, '1 nt rpretat ion of at hard so'lIar \-ray hurst"in Solar Activity and Re'la ted I nterplane tar5' a nd To _rost ridl Ph lhnomvlna,(Proc . First Europevan Astron. Meet ing, Athens, Sept . 4-9, 1972) (,id.J. Xanthakis) , VIl. I , p). 97 ('qpringer-Verlag, N.Y . ).
lloyng, P. , Brown, J..:, and van Beek, ll.F. : 1976, ''H-ighli tni rt sal tt iii iitil'sisof solar hard X-ray flairis obscrVyrd )in hoard Lte P51(t T'D-IA ;1( ~L' I i Io', Sol-ar1'hys. 48, 197.
14
Roferences
Kanie, S.R., Anderson, K.A., Evans, W.D., Klebesadel, R.W., and Laros, J.: 1979,"Observation of an impulsive solar X-ray burst from a coronal source",Astrophys. J. (Lett.), 233, L151.
Lin, R.P.: 1970, "The emission and propagation of -40 key solar flare electrons",Solar Phys., 12, 2b6.
Lin, R.P.: 1974, "Non relativistic solar electrons", Space Sci. Rev., 16, 189.
Lin, R.P., and Hudson, H.S.: 1971, "10-100 keV electron acceleration and emissionfrom solar flares", Solar Phys., 17, 412.
Lin, R.P., Mewaldt, R.A., Stone, E.C., Vogt, R.E., Paizis, C., and Van Ilollebeke,M.A.1.: 1978, "Spectral and temporal characteristics of solar electrons from20 keV to 20 MeV, B~ull. km. Phys. Soc., 23 4-AJ;.
Magun, A., Stewart, R.T., and Robinson, R.D.: 1975, "Three-frequency radiohelio-graph observations of solar continuum events", Proc. Astron. Soc. Aust., 2, 367.
M~tzler, C., Bai, T., Crannell, C.J., and Frost, K.J.: 1978, "Adiabatic heatingin impulsive solar flares", Astrophys. J., 223, 1058.
Nelson, G.3., and Stewart, R.T.: 1979, "Interpretation of a correlation betweenthe flux densities of extended hard X-ray and microwave solar bursts", Proc.Astron. Soc. Aust., 3, 392.
Ramaty, R., Colgate, S.A., Dolk, G.A., Hoyng, P., Knight, J.W., Lin, R.P.,Melrose, 0.8., Orrall, F., Paizis, C., Shapiro, P.R., Smith, ELF.,Van Hollebeke, M.: 1979, "Energetic particles in solar flares", in Solar Flares(ed. P. A. Sturrock) p. 117 (Colorado Associate University Press).
Richter, A.K., Van Hollebeke, >l.A.l., Hsich, K.C., Denskat, K.U., Keppel, E.,McDonald, F.ll., and Schwenn, R. : 1979, Max-Planck-Institut fijr AeronomicPreprint SP5-31.
Robinson, R.D., and Smerd, S.F.: 1975, "Solar flare continua at the metre wave-lengths", Proc. Astron. Soc. Aust., 2,374.
Schatzman, E.: 1961, "On the acceleration of particles in shock fronts", Ann.Aarpv., Lk, 2 3 4.
Smith, D.B.: L97s , "Mechanism~s for flash phase phenomena in solar flares", inCoronal Disturbances (Proc. IAU Symp. 57) (ed. G. Newkirk, Jr.) p. 253(Reidel , Dordrechit).
Smith, D.F. , and Lilliequist, C .G.: 1979, "Confinement of hot, hard X-ray pro-ducing electrons in solar flares", Astrophys. 1., 232, 582.
Stewart, R.T.: 1978, "Relationship between Type lll-IV radio and hard X-ray bursts",SolrPy. 58, 121.
Stewart, R.T. , and Nelson, 0.3I.: 1979, "An observed correlation between the fluxdensities of extended hard X-ray and microwave sol ar bursts", Proc. Astron. Soc.Aust., 3, 390.
Stewart, 8.1., Duncan, R.A., Suzuki, S., and Nelson, C1.1.: 1978a, "Observations
of high brightness temperatures In moving Thrpe IV solar radio bursts", Proc.Astron. Soc. Aust., 3, 247.
15
References
Stewart, R.'f. ,Ha nsen, R.'r. , and Sheridan, K.V . -. 1978b, ''Estimates of the- magnetic
energy densities of two eruptive prominences fromn their close association with
mioving Type IV radio bursts", in Physics of Solar Prominences (Proc. IAL' Colloq.44-) (eds. E. Jensen, P. Maltby and F.Q. Orrall), p. 315 (InSt. [hor. Astrophys.
Blindern, Oslo).
Suzuki, S., Stewart, R.T. , and Magun, A.: 1980 , "P'olarizat ion of herringbone
Structure in Typo 11 bursts", in Radio Physics of the Sun (Proc. tAU Symp.86) (eds. M.R. Kundu and I.E. Cergely) p. 241 (Reidel, Dordrecht).
Takakura, P.: 1972, 'The self-absorption of gyro-synchrotron emission in a magnetic
dipole field: microwave impulsive burst and hard X-ray buirst", SoiarI iyd., 26,151.
,,anl Beek, Ii. F., Iloyng, P., and Stevens, G;. A.: 1973, "Solar flares observe-d by thehard X-ray spec trometeor onl thle ES 'li ID-lA satellite' in Co 'cted Data Reports onthe August 1972 Solar-Terrestrial I'vents (od. Helen E.. Coffey). l'AC-29. Part T1,
p). U19. (1. S. Department of Commerce, NOAA, Boulder).
Van Hollebuke, >.A.I., Ma Sung, L.S., and McDonald, F.B.: 1975, "The variation of
solar proton energy spectra and size distribution with hl]iolongitude", SolarPhys., 41, 189.
Wild, 2 .P. , Murray, J1.D. , andl Rowe, W.C. : 1953, "Evidence of harmonics in the
spectrum of a solar radio outburst", Nature, 172, 5i3.
Discussion
Sch~wern. ll.re have boon recent direct meaIsuremients at I AU of thle elect ronscans ing Type Ill bursts, hiow do these elect ron numbers compare with thle bard
X-raV data?Stewart'. Lini and huidson (1971) estimated that 1;, of tile non -thermal ci cct rons in
tlt'' hard \ -ri\ source escape to interplanetary space. if the liard X-roes aredue to thermal CI'et rons tho numbehr (If lect tois escaping to space m~ight be dShigh as Ill.
Ccl fref cl: Were there any bursts whe~n a]lI four Possible MetLhods to es t imate thenumber of fast e lec trons were used?
St ewart : Lini and Hudson's (1971) list inclutdes several events where hard X-ray anld
rype I I 1-V bursts were observed to be associated with in tcrp Ianetar\ elect ron
events (see also Stewart 1978).
Alurka r: D)o Type 11I burst strue tures show a h ighi degree of p01 at i.at ion?
Stewart : fOnly tilie herringbone fear ures are polarized in the Cirllar sense. TheIordinary Trype I I burst with its narrow-frequencY bands showing slow dri it in
frequency with time, i.e. the ''backbone'' component of thet herringbone Typie I Iburst, does not appear tO be polarized.
C;iovanie li: At what level was thle l170 c field found Iv Olik?
Stewart: mil k 's model for thet imputlsivye nii eiwavc Source re-gion tad anl elet rondensity of 2 -10 Ill-% so it was piresumabl Iy nearlit thease oh tilit, corona.
Discussion
Croft: You mentioned that 108 K is needed to fit your data by invoking a thermalmechanism, but you added that no temperature so high has been measured. Yet your
phenomenon would last only about 30 s, would have a small area, and would occuronly rarely. I am not familiar with the operational history of those scientistswho measure solar temperature, but their methods might not work well in the
search for such a transient. Do you know if this is true? If so, then thelack of an existing observation doesn't necessarily preclude the thermalmechanism.
Stewart: I am not familiar with the technical problems involved in detecting lineemission at such high temperatures.
Note added later. Line observations of Fe XXV at 1.9 A during a solar flare
indicate a temperature -2-3 x 107 K. This is the highest temperature I haveseen quoted in the literature.
17
... ............... ... ... I I .
2. Magnetohydrodynamic (MHD)Description of Solar Radio Bursts
S. T. WuThe Universitv of Alabama in Huntsville
Huntsiife, Alabama 35807 U.S.A.
Abstract
The nature of solar radio bursts and their dynamic spectra usually arediscussed in terms of a static, semi-quiet corona. In particular, the value ofthe electron density and the general topological nature of the magnetic field,necessary for a proper description of the radio bursts, are deduced from un-disturbed coronal models. On the other hand we know that Type II and movingType IV bursts are related to flare-generated disturbances in the corona; as thedisturbances pass through the atmosphere, all plasma properties are changed.
In order to take these disturbed coronal conditions into consideration wehave used a two-dimensional MHD model to compute the properties of the disturbedcorona. Given these spatial and temporal properties, we have then deduced theresulting dynamic spectra of Type II and moving Type TV radio bursts. Threedifferent initial magnetic field topologies have been considered: closed-field,open-field, and coronal streamer configurations. While the model is unable tospecify the particular form of radiation mechanism present in any given
19
obs.rvat ion, z i, rtsu Its st rong lv suggest that t mpora I o I I - t 1 i D' l h,, hav i ormust he considered regardless of the mechanism (p lismni ,e,, i Il,it ins' r gvro-
synchrot ron em i ss i on)
1. INTRODUCTION
Generally, moving solar radio bursts ire usually believed to he caused
by one (or, in some cases, both) of the following nicilianisms: () svnchrot ronradiation from magnetized structures moving through the corona (see ftr example,Kundu, 1965; Wild and Smerd, 1972) which is called gvromagnetic emission andclassified as moving type IV, and (2) electron density enhancement in the ironawherein the plasma emission energy in Langmuir waves is converted into energyin escaping radiation either at the fundamental or the second harmonic of theplasma frequency. The evidence for fundamenta! plasma emission comes from harmonicstructure, first reported in Type II and Type IV bursts bv Wild, Sheridan andTrent (1959), and Wild, Smerd and Weiss (1963). '-merd and ),ilk W1471) note, inparticular, that the moving Type IV has three distintiv t'pes, namelv: (i)an expanding magnetic arch, (ii) an advancing front, and (iii) ;n isolatedplasmoid. In addition, they have suggested certain -oronal magnetic fieldconfigurations related to these moving tvpe IV radiio source-s as shon inFigure 1. Later, Kai (1978) has investigated further dotail I or the relationbetween circular polarization of the moving I'po 1V burst s and pluarity O thtphotospheric magnetic field. This stud% is based ,ii two-dimuis ional, high-
resolution observations (see, for example, Table I in Kai, 19781. other dataconcerning this subject have been given by Sheriden, Labrum and Pavton (1973),Schmahl (1973), Kat (1975), Stewart (1980), and Stewart et al. (1974). However,all of these works are based entirely on observational data. Smerd ind lilk(1971) have suggested that the second classification can b,0 explained bymagnetoihydrodvnamic (MHD) theory, but no s-st,'mit i, the,,rct i.1l udv dh.i [, 1.-nperformed for this or the other classificat ions.
In this study, we shall present an ideal, two-dimensional, t im-dependent,MIlD model which preserves the basic t oplogV suVggestid h% ill thr morphe logikalsuggestions made bv Smerd and Dulk (1971) as noted ,ibov. The important para-meters in this model are: (i) the initial topology of the magneti, t ield(Nakagawa, Wu, and Hlai, 1978; Wi et al., 1978); Iand (ii) A ub.littoisti, plasma? where is defined as the ratio of initial plasma pressure to magnetic
pressure at the lower boundary on the axis of svmmetry. In Section II, thebrief description of the governing equations will be given, and namer icalresults are presented in Section Il. We will discuss the results in thecontext of the observations, together with the concluding remarks, in SectionIV. It will be shown that certain observed features are reflected in thenumerical reF;ults. However, we will make it clear that the model is unableto specify the specific form of radiation mechanism (plasma oscillat ionsvis-a-vis gyro-synchrotron emission) present in any particular event.
20
2. BASIC THEORY
It has been recognized that the collective dynamical behavior of plasma andmagnetic field interaction can be described by magnetohydrodynamic (M1D) theory.The MHD, two-dimensional and time-dependent planar theory was first suggested byNakagawa, Wu and Han (1978) and Wu et al. (1978) to explain certain features ofcoronal transients which constitute certain observational manifestations of massnotion and shock waves resulting from a solar flare. It has been concluded thatthe density enhancement in the corona, leading to the observed white-light coronaltransient, is due to the interaction of fast- and slow-mode MHD waves. This modelis defined in the equatorial plane (governing equations are shown in Appendix I).Subsequently, Steinolfson and Wu (1980) and Steinolfson (1982) have extended thisplanar model to the meridional plane (governing equations shown in Appendix It) toexamine further characteristics of coronal transients. It should be emphasizedthat the equations differ only in the choice of the )lane of analysis and thatthe physical processes are the same; that is, the coronal response to a perturba-tion within it is one of collective ZMMD temporal and spatial reaction.
This model was compared with some success with the general characteristics ofobserved white-light coronal transients. More specifically, Dryer et al. (1979)have used an observed eruptive prominence as the forcing function in a model,which incorporated the observed electron density and ternrerature profile, toproduce some basic features of the white-light transient observed 1973 August 21,13:19 UT. Recently, Dryer and Maxwell (1979) made a comparison between radiodata and the fast-mode MHD shock wave computed from this model. In their study,the computed trajectory of the fast-mode MHD shock matched the trajectory of theshock wave inthe corona deduced from data on a radio Type II burst observed 1973September 5 for a flare which started at 18:26 UT. In this event, the forcingfunction was a 10 min. square-wave temperature pulse, the magnitude of which wasvaried on an iterative basis until the matched trajectory was accomplished.
Although this phenomenological simulation is not unambiguous, the final tempera-ture peak magnitude at the "flare" coronal base site was representative of thatmeasured in large flares. In this study, we will present computation results forplasma frequencies based on these suggested self-consistent MI) models with variousinitial magnetic field configurations as suggested by Smerd and Dulk (1971). Thenwe will utilize these results to interpret some observed radio signals.
3. COMPUTATIONAL RESULTS
In all of these computations, the steady state conditions for the system(both physical and mathematical aspects) must be established. Therefore, thesolar corona before the introduction of disturbances is chosen to be an atmospherein hydrostatic equilibrium with an isothermal temperature T = 1.5 x 106 'K, the0
number density at the bottom of the atmosphere, n = 2.7 x 108 particles cm-3;
and the poilytropic index y = 1.2. In the following, the specifics of thecomputational results are discussed.
21
:3.1 Closed Field C ase
As suggested bv Smerd and Dulk (1971), thle background magnet it- field con-figuration related to observed moving lype IV can be represented by three basictypes of tile magnetic field configuration as shown in Figure 1. We directattention, in this section, tothe first configuration depicted in Figure Ia. Wemay represent this configuration by (Wu et al., 1978)
Br= C r C2 cos 4
B: C 3 rC2sn~ I
vith C 3 =0.619 C, n C_) = 2.619, whe-re C I i, anl arbitrary constant to be
determined by thle choice of thle surface intensity oIf the magnetic fieldst rength.
W ith th is f iv Id -ont igurat ion and thle in i t ial atmospher ic cond it ionsmen t ioned earlI ic r, thle numer icalI computat ion i s sta rted wi thl thle int roduct ionot in sipots ive temperature increase of magnitude T =10 lo, in mid atmosphere
i.e., 20, 000 km above the lowor hou'Lndar') . Phvs icalIlv, LhiS enhancementcorresponds tt ;in impuls ive release of energy in a f lire, ;about 1029 ergs inthis calculation.
Figure 2 shows the contours of plaisma I fiindaim,,ntal I frequJelicy which changesbath temporal Iv and spat ia11v after introduction o 1 the' prototlpk "flare" Pulse.This calculat ion, made at t = 200, 400 and 000 sec( . was made for thle Case when
L!0=0.34, with .- being the initial ratio of te plaisma pressure to magnetic
pressure at the coronail base on the ax is of symniet rv. Obviously, , will initiallybe dif, nt at other locat ions; thus t he heti, as lus giv n , i s s implYchara, ist ic of t ie in it iaII it'lesent It ;tl 'If t le low corona. As timeinic teases, bet a will a I so Aiaiige i go ii I i it 1- eve r vwher. Al so-, some energy,resident init ial ly as magnet ic and t hermalI, will be converted to kinet ic ; hence,the Alfven Mach number will exceed one. In til; result, it is- clearlyindicated that plasmai tr co ve ni;io,:t l flis im .i loop-I ikestructurte which CX;iand ts I tI, i ut -IT l llti ; shis the initial
MIgtlilt i fie Illie I it di~hed 1I no) llld d11Illi muI nlgilet I I icid lines (Solid1lines ) bOO sec 1i't rlo duc ion m I t wit- 11- i Sohown 31rc thle local 'elOC i tVvec tors. Tliis f igir, corrt ponos t t 11C lami1rcqem'nlv enhancement contoursshiown in Figujre 2,. iCompar isn II t lit - two I iguirem Shows that tile loop-I ikevmissili struet lir, inl the lat t er t iglirc does, not conform idiretlIv to thle magnet icfjlux tubes1 sho0wn ill Figure, 3. Insteoii, thue (miss.i,,n structuire must, by;def nit ioi. -ntrm to, the- dc,itveiihneemen1.1I-t loop.
F igulre 4 shows,, iii i dilf fceit fiormat , tile plasma frequency as a function ofine at two, heights of 20,0100 km aind 35, 000 kil, respect ively. (Note that the
froquen,%v increases dowmwarl aloung the iirdinate of the figure. ) At a height of20, 000 km, thle plasma frequenc v sI owlv increases monot on icall v with t ime, how-ever, at 35, 0011 km. the p lasmij tIrequec, shows an inc rease in it l Ilv hut thena decrease after about 4110 sec. PhIVSiLicaIlv,. this difference in response at thetwe re-presentat iv ye iIt itudes may he interpreted in the foll owing wav. At thelower heights ( i .*., 21),0100 kin), whichi are liar thle e'Xplosive Site. beth tiliefast .ind slow MHDi waives iripigatc in it ia Ilv it almost thle same speed. Mathematically.tiis ciorrespotnds to a signal speed, C =v * ii (see Eq. (1 5) in Nakagawa , Wui andlian, 1978) . On thle ot her hand ;it the highier levels ( i.e. , 35,1)00 kin), the fast
22
°, S
#4 "
IM,
Figure 1. Schematic diagram of the magnetic field structure in three models ofmoving Type IV sources as suggested by Smerd and Dulk (1971). The magneticfield lines are shown as thick arrow lines, the paths of the moving IV sourcesbv broad arrows. The optical disk, the 80 Mtz plasma level and the averagestarting height of the moving IV sources are shown by thin arcs at R ,
Cee
.g R e and 2 , r aspectivel. The subscripts 1, 2 and 3 refer to the sourcepositions at three successive times (Smerd and Dulk, 1971).
and slow MHD waves are now propagating, at the later times, with their ownseparate characteristic speeds (see, Equations (12) and (13) in Nakagawa,Wu and Han, 1978). Therefore, the temporal variation at the lower altitudes ofplasma frequency are due to the interactions of fast and slow MD waves.
23
10) 200 sec Wb 400 sec
29 Ma149 MN,
9232 236166A
202225
0 1 0 1 2
110 600 sac5$ MHz
CLOSED FIELD CONFIGURATION
1,0 -03
244
0
Figure 2. The contours of plasma frequency enhancements at time 200s, 400s, and
600s after the introduction of the disturbance for 0.34 and closed field
configuration.
Figure i. \'lIoc ity f ield and magnet ic I eld Wit igllrat ions 000ts after int roduct ion
axis (vertical In this figure) ,I, t.,- thi Ii.j 1. J l In il igtir' 2(WuJ "1; 1978).
24
too
CLOSED FIELD CONFIGURATION
x 9.034
0 SO
too- $0.000 k.
02 500, 400 -; DO go 00o
TIME (ee
Figure 4. Plasma frequency versus time at heights of 20,000 km and 35,000 kmfor o= 0.34 and closed field configuration.
3.2 Open Field Case
We now turn our attention to the second magnetic field configurationsuggested by Smerd and Dulk (1971), as shown in Figure lb. This topology,essentially an open (radial) field configuration, may be represented by thefollowing representative equations:
B r = C, rC 2
sin 1,
B = -C r-C2 cos , (2)
where C1, C2 and C3 are constants defined as shown previously in Eq. (1).
With this initial magnetic field configuration, together with an atmosphere inhydrostatic equilibrium as mentioned previously, the numerical computation againstarted with the introduction of an impulsive temperature pulse (as in the case ofthe closed field topology discussed before). The contours of plasma frequency foro = 0.34 are plotted in Figure 5a, b, c, respectively, at t = 200, w:,0 d 600seconds after introduction of the pulse. The corresponding magnetic and velocitytields are shown in Figure 6, in which the dashed curves represent the magneticfield lines at t = 0, and the solid curves indicate the magnetic field lines att = 600 sec. From these results, we note that the "prototype" emission coming fromthe fundamental (or even second harmonic) plasma frequency has formed as a plasmacloud-like structure preceeded by a quasi-parallel shock. This is in contrast tothe loop-like structure appearing for the closed-field case. Also, we note thatthe field has been pushed sideways by this enhanced plasma cloud ahead of whichthe quasi-parallel shock has formed. Figure 7 shows the plasma frequency as afunction of time for 6 = 0.34 at the two different altitudes which were considered
0previously for the closed-field topology. It shows that the emission source has aminimum frequency at 20,000 km. At 34,000 km, the emission is essentially constantfor 400 sec, after which it rapidly increases in frequency to - 300 MHz at t =600 sec. Eventually, however, this will return to the background level after thedisturbance passes through.
25
(a) 200 sec () 400 sec
153 MNz , 1-53 M11
258 361269 391
0 2 0 I 2
Wc 600 sac(c 60OPEN FIELD CONFIGURATlON
153 MHz 0 34
215
5119
0 12
Figure 5. The contours of plasma frequency enhancement (solid line) and depletion(dashed line) at time 200s, 400s, and 600s after introduction of disturbance forso
= 0.34 and open field configuration.
\ '\
Figure 6. Velocity field and magnetic field configuration 600s after introductionof the disturbance for open field configuration (Wu et al., 1978). Note again,that the equatorial axis (vertical In this figure) corresponds to the horizontalaxis in Figure j.
26
OPEN FIELD CONFIGURATION
3 200 20.000 -
220
I 3000
350-
40000 0 400 g00 S0
TIME (ee)
Figure 7. Plasma frequency versus time at heights of 20,000 km and 35,000 km for0.34 and open field configuration.
3.3 Coronal Streamer Case
According to the suggestion of Kai (1975), some radio sources were observed ina region which apparently consisted of a background magnetic field which had a con-figuration of a coronal streamer as shown in Figure 8. In this section, we willdiscuss the temporal and spatial variations of the plasma and magnetic fieldresults, obtained from precisely such a configuration, in more detail. The coronalstreamer configuration is obtained by solving the set of governing equations shown
in Appendix 11 together with a predetermined solar wind. The detailed account ofthis work can be found in Steinolfson and Wu (1980) and Steinolfson, Suess and WU(1982). A typical solution of the coronal streamer $o = 0.5 (defined in the
0
,ame way as betore) is shown in Figure 9. Shown in this figure are the following:(a) magnetic ficid lines, kb) pressure distribution, (c) velocity vectors and (d)density contours. The vertical axis is the equator and the horizontal axis is thepole, both extending from I to 5 R I( being solar radius). The short dashed linein Figure 9(a) represents the Alfven mach number (u/V ) being unity, where u is
12 Athe flow velocity, (V +V )11
2, and V is the Alfven velocity - . The longer
r 11Adash d line represents the sonic Mach number (u/ai being unit', where a is thesonic velocity (.R1')'2. It is seen that a steady-state, self-consistent solutionfor both a coronal (or helmet) streamer, together with a polar coronal hole, i,provided.
With thLis solution as our initial condition, the "prototype flare" computationwas carried out hv introduction of a thermal pressure pulse at the lower boundaryof the corona 1i.e., R, -I). This thermal pressure pulse has a magnitude of 10 inwhich the temperature is increased by a factor of 5, and the density hy a factorof 2. I'his pulse is kept "on" during the entire computation. In Figures 10 and 11,we show the coronal magnetic field and density responses at 80 and 180 minutes afterIntroduction of the pulse. In these figures several interesting features are worthyof ment ion:
(i) In the closed magnetic field region, the field has been pushed up bythis pule, thereby enabling its interpretation as a rising arch;
(ii) This closed magnetic field region eventually is pushed open at latertimes (see Figure 11) while the plasma pressure and kinetic energy density becomecreater than the magnetic pressure; thi condition will relax back to the sun
27
i P~ellu L- d F-'c
"aI
(d)
Figure 8. A schematic model showing the successive appearance of a I are-continuum source, an isolated moving sot rce and a storm-continuum source withbi-polar structure (Kai, 1975).
following cessation of the pulse;(iii) A MHD fast shock is formed and propagates outward;(iv) A high density region forms behind the shock as suggested by the
speculation of Smerd and Dulk (1971) shown in Figure 1. This high density regionforms a loop-like structure; however, the maximum density enhancement (i.e., plasmafrequency enhancement) is immediately behind the normal shock, and decreases some-what along the optical, loop-like transient's legs. This phenomena shows features
similar to the model (see Figure 8c) suggested by Kai (1975) based, phenomenologi-cally, on the observations.
(v) At a later time (i.e., 180 minutes after introduction of the pulse), mostof the original closed field region has become open, and the center part of thedensity enhancement has propagatod out of the computational field of view (i.e.,beyond 5 Rs). The density-enhanced leg (within the two short-dashed lines) remainsas shown in Figure 11.
28
CORONAL STREAMER 0-.16 hrs. P-0.5)
S (~~.) MAGNETIC IPMSFIELD LINES
4--4
4 5
* .(4) VOLOCITY S(4) DENSITYVIECTOO4S
lo., W~
Fi;ure 9. Planar maps of the~ magnet ic field cont igurat ion anu plasma propertiles inthe coronal st reame r for O=0. 5 (Steinoltson and 1,o , 1980).
(t=80 mmi, eo=05)
a 0 -
-- ,, rE.
a..:"
F-igure 10. The magnetic field configuration and densitv distribuition 80 (ninu, s
afteor introduct ion of the di sturbance for e, = 0. a The lonyg 1all-.hd I 1i1V
(a) MAGNETIC W UENS TYFCO LINES
4 / 4
r.2 I
:11.
I 2 3 4 5 21
Iter i I,,tr billet in it tl 1 st ir I, 1I 1, 1 1 1, ther svmbills art, the'mc is i..~ure 10 (Ste-[0 l t ,I Ojl 1[,1 !- I18I)1
t. CONCLIIDING. REM1ARKS
III this Itiiilv, W" had :1r, Sclit_-4 j ---. imI), :"11i, 'In-b', : ,
ilinputat on for thr-' t'.pv S] 'I nglt 1 1.1> iiiri ii] h.have heen -iuggcstedi hi ,I number .1 rid i.. i~t r,.n..: t - X; 1 il 1 rcv ?.I .s
oi r ols1 reIt ii.- 1, u halve .t aIrt ed I, i t I k -:I ;, tI, dI rt
tiis (Ii s tirhsiic mii' I, i It .rpl- I~ .I I .1 , ,rI t ,I- Ir ac in d , r sLti %aic t Iv it v Vw1 tilT . cvlilt il )t tls h, 1im I t t17" U-
\e jIoc I tv Id Tiiglc I, i c I d Ii ill I l 111 I II's
l)JIri 'te cr., c 1 t n t% ia d ill L. h I t i,-J 1,1 ;
s erl Ila sses e I rIidi,. sor( ,s ire 1) rL.1T I:. h1 r rr d I, 1 ImII~Le ,]Il1) 1 asmai !requenc,l ou gh llr r-suits, 'i itteIr7, :,: li~i m z ll bIT>
preIl' I~l 111115T iii..r t. ilrl 1' tihe ve slid ? 'r : I Itl r I i i rm I
Ililract er i st I s -nil ht su ii;r i /t-d is t , I IeWs( i ) -ronm thIiis ,ii I Iit ionv , e Ic I , r i h11 1 t l ie h I ! nI tI l : I
desnlihed Ilere i aS al dens itv or, equi ialett I %, plasma I rei:toencv enliai vtI-rilform-; at 1LlI111-1 ike s t ni ',lii] Whi Ii prlil u sIttward a. s 1wwr ii Is F I
hehav Ilolr has t he' halic a1 t I.r istL I of r is ins,, arch . For the open I itl I 'Illi it ion
cillplitLation , theL 01 ei l il oree re sit t as i niih I c- i ike 1.1 i 1 1 ii h! I101-
gates outward, i5 shiownl ill Figure )1. Fiirt Itier , - 'Tn, 1 11 ~tl " that !5 h III til h, ; I
Vil id (olr dens it v dep It vLin rg il) I Sh IOW]] . IllS2 CI I Iltl tedr liIit lYt. b1 .~mI theIIclse sugge-stetd hb. Sme rd aodl DI k (I1 971 Iiledil l ohserV,It il A, 'hulwv :11 F igliri 1 l
Fi nalI lv, th [iI ,rilI st rehIther 1 a.1CUI l t io (11Show 11 thit (CM eti ssili 5.inl. so li-hvetwil-flild chariterist ic. ('art It thek silliri-es prop~agaite iiitWAtil, W1ti Iteh' tiletrpliottion remilts bhlind w i titn ti ItI g reg iri Il''5 t Il til e -t (I 1 'ii I
hii I i7 rIt ) Is sho nI Ii I ii
301
I F- 7
(ii) By presenting the computed results of plasma frequency versus time atseveral representative altitudes, we show how the plasma frequency (fundamental or,equivalently, its harmonics) behaves after the disturbance passes through the corona.These results demonstrate, once again, that the magnetic field configuration in thepre-flare state plays an important role. For example, the closed-field case showsa U-shape behavior at 35,000 km; whereas, the open-field case shows an inverted U-shape behavior at the same altitude. (Note, here we used the "U-shape" in a frequency-time context which does not imply U-shape intensity burst.) A word of caution con-cerning further interpretation of our model is warranted. Because of the limitationof the one-fluid IHD theory, we have no way to discriminate between mechanisms ofvarious radio emission; that is, we cannot specify whether plasma oscillation or gyro-synchrotron emission is responsible for a given case. However, some gross behavior ofradio emission, using the plasma oscillations as a working hypothesis, has been,emonstrated through this type of calculation. Gyro-cyclotron frequency (or itsharmonics) could equally be employed by referring to the temporal and spatialmagnetic fields. Further study should include this latter point at well as a modelin the form of multi-fluid MID theory in which the detailed gross dynamical inter-action between ions and electrons can be studied. In this latter case, polarizationeffects can be shown. Thus, some detailed characteristics of radio emission sourcesresulting from theoretical calculations such as that discussed here may be compareddirectly with observations.
Acknowledgments
The author is indebted to fruitful discussions with Drs. M. Drver, R. Stewart,K. Sheriden, Y. Nakagawa and F. Tandberg-Hanssen. In particular, Dr. M. Dryer readthe manuscript and gave many valuable suggestions; Dr. R. Stewart brought to theauthor's attention many valuable references when the author visited CSIRO, Divisionof Radiophvsic, in Sydney. This work was supported in part by NSF Grant ATM77-22484 and in part by NASA/MSFC Contract NASB-28097.
31
Appcndix I
The basic ideal MHl) (without dissipation except at theshock) equations fornumerical computation are written in the following form in spherical coordinates(r, 0, 4) and specialized for the equatorial plane (0 900).
Mass conservation:
, +1 (2 r)
it+ rr +,-V) + (,,V,) 0 (At-l)
it 2 r rr
Momentum const-rvation:
V r V V 2,
t r )r r r
+ B; + B il. .. 1 .. . . B - (AI-2)
r r r r 2
V \ I. v ,-J- + V _ + .- ..
t r r I r
-r B B
--. -r + - - (l-r r r
Induct ion equat ions
B
[r
Br =_ [. r(B V -B V ), (Ai-4)
2 r r r
+ [2 r { rr + - I Vr + V2), r -- r
+ r 2 2(Vr - Virli;
(Equation (AI-) continued on next page.)
32
+ -- V p + 1 r(V +V 2r i'p YPu- r p
r 2+ (V B - VrBrB)
+ Vr "GMs 0 (AI-6)
Here, the dependent variables are the density ;,radial velocity Vr , azimuthal
velocity V,, thermal pressure p, radial magnetic field Br, and azimuthal magnetic
field B P; the constants are the polytropic index f, the gravitational constant G,
and the solar mass M . In addition, the equation of state is used, namely,s
p RT, where T is the temperature and R is the gas constant.
Appendix II
The two-dimensional, time-dependent ideal MHD equations in the meridionalplane are written in the following form:
Mass conservation:
2;V V~co i,(AII-l)
+ 3r ( + rr = r r Cot
Momentum conservation:
11' V N V V2
r + V r + 6 r V't r r r .-A r
2B 'B B 'B B2 GM
I Pr 0 6 a sr_ + r je r w 2 (All-2)
r
'V H V Ve , Ve V rV 1-+ V +---
t r r r r
B 3BE B 3 B B
= _r _mi 'p r rr 6; + wior (AII-3)
33
Induct ion equations:
,B V B -V B_r _- r .. .. r 1 ( r V Bt r (V V. .) cot l(AI -4)
B, (r -V Br) = (V B -V B), (All-I)
r r ( r ': r
ELne rgv colns,.vrat ion
'V V V,-_P + _r + V 'P + :y_ --- + _T' ( II 6
't r * r r rA 0
- ~2% + V cot ). (A I 1 -6)
lere, all symlols h1yV, I , C .l;7. ",!cniI;2 7§ , . , I i.t I iii lIlt' . Ia d d i t i o n I , t V i d t I 'p l ' 'I t'H 1 t : ' : .1 1d i o bl I' s; I, i td ai n dl I % ' 1.1 1 1 ,1 i i i , ! la I L i i
f ie ld, rI'' C LivvvI v,
References
Dryer, '1. and M'axwell, A.: 1979, "Radio data and a the,,retical nodel for thefast-mode :1111) hock wave generated by the :,;olar flare of 1973 Septerher 5,1-3:26 ITT", A,;tr,,j)J17,;;. 7., 2?1, 945.
)ryer, M., Wu, S. T., Stetnolfson, R. S., and Wilson, R. A.: 1979,.agnetohydroJynanic riodels of coronal transient, in the meridional plane.I1. Simulation of the coronal transient of 1973 August 21", Ast Orph'Js. I.,227, 1059.
Kat, Kelzo: 1975, "A nodel for the developnent of a solar outburst based oMobservations with the Culgoora radio spectograph and heliograph", S;o.r Phys.,45, 217.
Kat, Keizo: 1973, "Relation between circular polarization of moving Type IV hbirstsaid polarity of photospheric magnetic fields", S.'lir Phy;s., 5A, 417.
Kundi, M. R.: 1965, So!iar Radio Astroniomy, Intersctence, New York.
Nakag~wa, Y., Wi, S. T., and lian, S. M.: 1978, "Magnoetohydrodynamics of atnoqpher,-tran-;Ients. T. B;asic results of two-dtinnslonal plane analyses", At rphi. .7.,21), 114.
34
References
Nakagawa, Y., Wu, S. T., and Ilan, S. M.: 1981, "Magnetohydrodynamics of atmospheric
transients. 11. Basic results of non plane two-dimensional analysis",Astrophys. 7., 244, 331.
Schmahl, E. J.: 1973, "Catalogue of nine moving Type IV radio sources observed withthe Culgoora radioheliograph", Australian J. Phys., Astrophys. Suppl., 29, 1.
Sheridan, K. V., Labrum, N. R., and Payten, W. J.: 1973, "Three-frequency operationof the Culgoora radioheliograph", Proc. Inst. Elec. Electron. Engrs., 61, 1312.
Smerd, S. F. and Dulk, C. A.: 1971, "80 Mhz radioheliograph evidence on noviagType IV bursts and coronal magnetic fields", in R. Howard (ed.), Solar Magnetic
Fields, IAU Symp. 43, D. Reidel, Dordrecht, 616-641.
Steinolfson, R. S. and Wu, S. T.: 1980, "Dynamic simulation of coronal massejections", UAH Research Report No. 240, University of Alabama in Huntsville,Huntsville, Alabama.
Stetnolfson, R. S.: 1982, "Coronal loop transients in streamer configurations",
Astron. and Astro. (in press).
Ste[nolfson, R. S., Suess, S. T., and Wu, S. T.: 1982, "The steady global corona",Astrophys. 1., 255, 730.
Stewart, R. T.: 1980, -Transient disturbances of the outer corona", in M. Dryer and
F. Tandberg-Hanssen (eds.), Sl~r and rnterplanetary Dynamics, lAU Symp. 91,). Reidel, Dordrecht, 333.
Stewart, R. T., McCabe, M. K., Koomen, M. I., Hansen, R. T., and Dulk, G. A.: 1974,
"Observations of coronal disturbances from I to 9 RO, Solar Phys., 36, 203.
dild, 3. P. and S.ord, S. F.: 1972, "Radio bursts from the solar corona", Ann. Rev.
Astron. Astrophys., 1,1, 159.
Wild, J. P., Sheridan, K. V., and Trent, G. A.: 1959, 'The transverse motions of the
sources of solar raiio bursts", in R. N. Bracewell (ed.), Paris Symp. Radio
Astron., Stanford Vniversity Press, p. 176.
Wild, J. P., Smerd, S. F., and Weiss, A. A.: 1963, "Solar bursts", Ann. Rev. Astron.
Astrophys., i, 291.
Wu, S. T., Dryer, M., Nakagawa, Y., and Han, S. M.: 1978, "Magnetohydrodynamics ofatmospheric transients. II. Two-dimensional numerical results for a model solar
corona", Astrophys. J., 219, 325.
Discussion
Gelfreich: How could you get twisted magnetic fields? Is it due to initial con-ditions or is it a solution you have obtained?
wu: In this calculation, the twisted magnetic field is caused by a rotationaldiscontinuity. Details are shown in a paper by Nakagawa, Wu and Han (1981) who
show the twist in a self-consistent MMD model which considers the tme-dependent
evolution of all three components of the magnetic topology and plasma velocity.Yes, the initial field is being twisted.
35
Discussion
Zintni k: You hel love gvrurctsonanCe emlissin to he responsihle for Type IV bujrts.What about electron velocities: are the electrons relativistic or nut? And howcall vou explain fine stroicture on Type TV bursts In the framework of gyro-frequency radio emission?
Wa : In this simple fluid 11I1D model, we cAnnot explain those fine structures; however,we plan to move ihead to extend this model to two fluids and non-plane analyses.Then, it should he able to denonstrate, in pirt, what you are asking. That is,we expect to -,hfow the local ized regions .- f local ly-compressed magnetic fieldswhich, lin principle, could - together with the few electrons in our (assumed)Maxwellian distribution function - conspire to produce gyroresonance emission.
CI4t III tic'orc-t ic-u st tiic-, sillil -Il 0111 lMiut tI1111OiMate' inl ordecr toma kc t ,- rc A II,,- '.1 t I I. raictcdal- 1, i 1, fit ti, clit . I i- fit in d ing which kinds ofSimpli i cii icii In c ! id, wit' .t th 1r & -1'I i it illt.' -cntof thet moreImportant phi,-ctci:.t-iic. tilt'. ,-r.t ic -), t t I.tx r oI - !!t i , tc - fo rm calIculat ions(as , oct hai,ivc and I t, ic-t to look to ii sii' la t i t',,1 1 ii . t [It- dCt'i Ve'd re.sults andthe corres, poni iti *h%ca c i 'elits xpyiuci inmcntal I.,icevd How well1 dci yourmethod, ic-rtorm, Wh.l itt nh d b''tis apprOat'
I ut , h i 't 'ir -iOS, :- I i ' I , i tI II c, ticc''[, -,, ioci-, I v-pub] i shed rest.- t ocif
a se r ies o1 t a-p,-i.,h i oi; ei tci: I s 1 LI LL c. it) Ofcsci u these papers(plea,-' -eo tlie rc, i~,: lisct to the, wtittcn x'-sOn this talk as wellas my r epl1V to Or. . icc tin)w demeust rat.': (a) hoc~k dc-i~tcttand' vc'loci t
as measured 't'.* Vp LVic al lx pc ii rad in hUrSt s ( h) dcis it'. enhancements, foallowed
by rare fact ions,* ci- mc as-icd !iv coronograph obscrV',tiLons o t K corona scat teringof suI i gilt tromn tratens it I v-d iustur!hed Icct rons; aind (c) shiape of coronaltransients as alIso clctc' u" ned for a number ofci as- I'% Lt co ranagraph.
S, itwei;l I~c It, cm inI tes,, c aIcu lat i on - k'a (ils tc ;IloalI t, rip- ra turc I n Creast1) id Voilt Icyt c on., i d, r, a 1ii 11 tc rta t iv,',, t ie, i n c-ct , uin oi ttc',w miatcc r ialI
Wu Yes, %we havc'- coens iderci various a ItLt t rat iVCeS. I rca' r-c t %Oxil to Lte paplerbv Ste i ncilson ct ail. , to appear in the. P roceed ingsfth ugtasncnene
'lu .ini 1%. In that rpter we used emc.rginci ma'nc't ic flux a,. our
d istutlitc. II sevtral o thter papers in t he Atrpisic 'Icra we used vari -
ous combintationts ct enhanced dens itv (to simulate i 11 ecat ici iii new plasma t rom
tice e hrcimcspherc' into Lithcase of tiiecircita) and tempera tare pclses. In tune Of
Lt-sc i 1~u , 1;)[,,: et a!., t.,rojlvst .j_,22;7, 1cm1') I- 1lcc'Ji IlOn4-lai-tng
dens itv and tCHIiieraturc' enliancementL as measured 1)% tlc Nc\Sc/MISP(-Aerospacc'soft X- ray inst rument aboard Skvla h ug thes 21 Aiigci-t 1973 1limb event.fiII, t.il t imult-I it i ditt fcatic7 (It thel lIberVeC coronal. transient asobctcl
I,% tle NC(ARIiAii eoroncgraphh.
I .i ;'. no; I t i u,( i :i citII t tc oc s c.v 1ow s uc cc t eI I svt'mr c ou Id dl.uIc-'xlcc. g ive,,nthat thIc' lii- 'sIf to re- all orig inati fro0M ti0 Magic't ic (ictic's in tile model used;or thiat , as itt Vcu' s case, -itch a model f ield could develop arc c'plos ivyC
instabti it-' IcIcceet' a S ipjlc chiatg' fin ciodc'l could hc' cv ised to provid,- an
instai ii t w i it 1 ,it ciil ,,i i t~i i -aI I t chan111ge C i L it D I I i In at l"x' i s 0t 01C 1!I%.' ic L
Wit; It i', --c~ litial to orttct t a mimsintc'tptc'tatitin lit I iiav iavc''-nintc-t ntialic,letft with Ii 'uc (.i itxplic l iii ''1 r rc'ma tk ). lte dc nlot ci clnc thtat atn ex;ulcs ivinstabi litv 'cctiIs inii ititI c(511iuirium (I tC., )antic-Ilt ) Magnetic f-icid andatmoshirc- . icit hitt, llc' , imltIatc' ,ian cnctg ,cinvt'i S ictil prtc'' (as . tor c xampl,by rt tciic I ini II. varimi i x i I ils thIr o 1ii tiict -i ii cc ottsc'l c it1 c- - ait i Sl it(
pillse (wit itir ticul,lvaI imc tempocracl profc i Ic-s ) r a magnct it' pitlse Ice. g.
atn em,tii,,,itn, :itl. ic~ t, it I I . it- I it-st twot Cast's Wct hiuv shotwn may not fit'real ist it-; litIw,'vc-i t Ic l ii do -\hiii t fiittimnta cii t i c milt i third ca--
(I. c-. , tt I . t ,-r i tIt gi-tit ln Iici t c.'ci:1i1,l c. Ilt' I I v iicr Va t tuna I
anaIlv, i.
3. Profiles of Coronal MHDDisturbance Associated With
the Solar Flare on April 27, 1979
R. V. Bhonsle, S. S. Deaonkar and S. K. AlurkarPhysical Research Laboratory
Ahmedabad -380009, India
Abstract
A solar flare of importance IB which occurred at 0636 UT on April 27, 1979 inregion 1705 (McMath No. 15967) on the solar disk (N20, E16) produced intense radiobursts. An impulsive microwave event was also recorded at 2800 MHz at Ahmedabad,beginning at 0642 UT, and showing about 10-12 peaks in intensity in a time durationof 20 minutes. The most interesting feature of this event is the observation of astrong continuum radiation (Type IV) starting at 0653 UT and lasting for about 10minutes as recorded by the wide band spectrograph operating in the 45 - 25 MHzrange. This continuum radiation displayed sharp low frequency cut-off, whichvaried from about 40 to 30 MHz in a quasi-periodic manner. The sharp low frequencycut-off In the continuum radiation is attributed to *Razin effect". The perturba-tion of this cut-off frequency is interpreted as that induced by the passing MHDshock wave through the magnetic flux tube whe'e the continuum radiation is gene-rated by the synchrotron mechanism from trapped energetic electrons. Assuming aspherically symmetric electron density model and using the observed cut-off fre-quency, the magnitude of the coronal magnetic field around 2 R. above the photo-sphere had a value of about 6 gauss.
37
1. INTRODUCTION
A wide band decametric spectrograph with medium resolution (- 50 kHz recently
began operating at Ahmedabad, India (230 02'11, 720 36'E) for the purpose of recor-ding burst spectra in the range 45 - 25 MHz. As the radio bursts in this frequencyrange originate from the coronal region I - 2 solar radii above the photospherewhere transition from closed magnetic field lines to open field lines is believedto take place, it is interesting to deduce properties of the region from thecharacteristics of the burst spectra just as the microscopic features of coronal
density irregularities were deduced from high resolution decametric spectroscopy(Sawant et al., 1976; Bhonsle et al., 1979). In the present paper, we report onan interesting solar event which was recorded on April 27, 1979 with our wide band
spectrograph.
2. OBSERVATION
A sotlar flare of optical importance lB and type X-1 in X-rays occurred at
0652 UT on April 27, 1979 in region 1705 (McMath No. 15967, location N18, E17).Though the Boulder Solar and Geophysical Data reports (1979) indicated that theflare was not spectacular optically it produced intense Type IV radio emissionwhich lasted for about 10 minutes starting from 0653 UT. Figure 1 shows the first
minute of the Type IV burst spectral record which reveals some interesting fea-tures. Following th. solar flare at 0652 UT, no conventional Type II burst wasseen on the spectrograph record but a strong Type IV continuum radiation was evi-
dent from 065] UT t,) 070Z UT.
It can he seen from Figure I that the Type IV radiation appears with a sudden
large decrease In cut-off frequency from about 44 MHz in a time duration of about8 seconds beginning 0653 UT. Then the cut-off frequency increased to 42 MHz afterwhich the radiation spread again slowly to the low frequency side to about 33 MHz.
In addition it can he seen that the observed cut-off frequency of Type IV emis-sion, while it was slowly decreasing, underwent a quasi-periodic perturbation with
small amplitude. This is seen more clearly in Figure 2 which is the tracing ofFigure 1. Thus, the initial sudden change in cut-off frequency is of the order of9 MHz and the average frequency change and period of the perturbation as determined
from the record are + 0.3 MHz and 5 qeconds respectively.
3. INTERPRETATION XND DISCUSSION
The perturbation seen on the low frequency side of the Type IV continuumappears to be the result of the dynamic interaction between the magnetic field andthe material ejected during a solar event. Numerical solution of the time depen-dent magneto-hydrodynamic (MIlD) equations of motions have been obtained in somesimple situations to understand how a disturbance propagates upward and laterallythrough the corona with the leading edge of the disturbance having the shape of anexpanding loop (Steinolfson et al., 1978). The leading portion of the disturbancecontains only coronal material whose properties have been altered by the preceding
2.
38
AHMEDABAD
APRIL 27, 1979MHz WIDE BAND SFECTROGRAPH
24-8
34.2-
06300653.5 0654.0 06541
TIME U.T.Figure 1. Dynamic spectrum of Type IV burst observe~d on April 27,1979 starting at 0653 UT. Note sharp low frequency cut-onff ofcontinuum radiation and its quasi-periodic variation.
AH MEDA BAD
MHz APRIL 27,19?924.8 WIDE BAND SPECTROGRAPH
24.8P
34.2
4 3.5
TIME,0 UT3. 0654.0 0654.1
Figure 2. Sketch of hurst spectrum shown in Fgure 1.3.
39
waves (which may strengthen into shocks) created by the explosive nature of thesolar event. Hence we may assume that the perturbations seen in Figures I and 2,
caused by the MHD disturbance along the closed field lines forming the magneticbottle in which energetic electrons are trapped, give rise to Type IV emission bysynchrotron processes (Takakura, 1966; Wild and Smerd, 1972). Such perturbationscan be expected to produce magnetic field fluctuations which in turn modulate thecut-off frequency of Type IV continuum (Ramaty and Lingenfelter, 1967).
Ramaty and Lingenfelter (1967) have shown that the sharp low frequency cut-offJn Type IV continuum may be explained as due to Razin effect rather than on thebasis of cut-off imposed by the local plasma frequency. The latter possibility
does not arise in the present situation because the generating mechanism of Type IVbursts is a synchrotron process which is independent of the local plasma frequency.The Razin effect is seen only when the parameter a = 466B/Ne
l/2) is less than 1.
Itere B is the magnetic field in gauss and Ne the number of electrons per cm13
in tilesource region.
Newkirk (1971) and Dulk and ,IcLean (1978) have reviewed the coronal magneticfield estimates. However, we can estimate the magnetic field value in the coronafrom the observed Razin cut-off frequency from the relation given by Ramaty and
Lingenfelter (1967).
Neu R - L"0 1 (1)
where il R is the low frequency cut-off of the continuum in Hz. Putting ) R
- 35 x 106
Hz and He - 107 per cm
3 at 2R., using Newkirk's model (1961), we find
that B has a value of - 6 gauss which is in agreement with the estimates given hy:;ewkirk (1971) and Polk and McLean (1978). The initial sudden impulsive change i-lcut-off frequency at 0653 UT of about 9 MHz and the average amplitude of perturba-tion of the cut-off frequency of + 0.3 MHz makes the percentage change of - 30 and- 1 at the observed cut-off freqency (- 30 MHz). Hence, the change from I to 3f.
in the magnetic field, B, is due to a MHD wave at that level in the corona.
Another interesting feature seen from Figures I and 2 is that four "puffs" ofradiation appear to emerge when the cut-off frequency of Type IV radiation reachesminimum. This shows that energetic electrons are ejected in bunches from the mag-netic bottle when the magnetic field is quasi-periodically strengthened. Thereleased electrons subsequently caused radiation by a plasma oscillation process.
4. CONCLUSION
From the observed low frequency cut-off of Type IV continuum burst in the de-
cameter range, we have obtained estimates of the coronal magnetic field (- 6 gauss)around 2 solar radii and the magnitude of its perturbation (from I to 30%) as aresult of an *IHD wave propagating through that region.
40
Acknowledgments
Our thanks are due to Prof. D. Lal, F.R.S., Director of PRL for his interestin this work. The cooperation of the field staff in the maintenance of the ,'quip-ment is very much appreciated. This program is supported by the Department ofSpace, Government of India.
References
Bhonsle, R.V., Sawant, H.S., and Degaonkar S.S.: 1979, "Exploration of the sclarcorona by high resolution solar decametric observations", Space Sci. Rev.,24, 259.
Dulk, G.A., and McLean, D.C.: 1978, "Coronal magnetic fields", Solar Phy., '7,279.
Newkirk, G.: 1961, "The solar corona in active regions and the thermal origin ofthe slowly varying component of solar radio radiation", Astrophys. 7., 133,983.
'4ewkirk, G.: 1971, "Large scale solar magnetic fields and their consequences",in R. Ioward (ed.) Solar Magnetic Fields, Proc. lAU Syrp. No. 43, D. Reidel,Holland, p. 547.
Ramaty, R., and Lingerfelter, R.E.: 1967, "The influence of the ionized medium onsynchrotron emission spectra in the solar corona", J. leopliys. Res., 72, 879.
Sawant, H.S., Bhonsle, R.V., and Alurkar, S.K.: 1976, "Microscopic spectralfeatures in solar decametric burst and coronal irregularities", Solar Phtis.,50, 481.
Solar-Geophysical Data: 1979, U.S. Department of Commerce, NOAA, Boulder,Colorado, 417-1, p. 7.
Steinolfson, R.S., Wu, S.T., Dryer, M., and Tandberg--lanssen, E.: 1978, "Magneto-hydrodynamic models of coronal transients in the meridional plane. I. Theeffect of the magnetic field", Astrophys. 7., 225, 259.
Takakura, T.: 1966, "Implications of solar radio bursts for the study of the solarcorona", Space Sci. Rev., 5, 80.
Wild, J.P., and Smerd, S.F.: 1972, "Radio bursts from the solar corona", Ann. Rev.Astron. Astrophys., 10, 159.
41
Discussion
Dryer: Uhat is the reason for the cutoffs which gave the results of nultiplc"puffs" at the lower frequency?
\.Arkar: This may be due to the energetic electrons released in bunches from themagnetic field bottle when the latter is quasi-periodically perturbed by theM.H.D. disturbance.
Celfreich: Was there any evidence of a Type II burst in this event?Alurkar: No, there was no evidence of a Type 11 burst.
Wu: Jhile you determine the magnetic field strength from the observedfR 20 .18, you need the disturbed number density. Haw do you get thisdisturbed nmber density?
Al urkar: Newklrk's values of '4e give a rough estimate during undisturbedconditions. For more precise values number density in the source region shouldbe used.
42
00
0\
0
4. Fine Structure in Metre-WaveSolar Radio Bursts
D. McConnell and G. R. A. EllisDepartment of Physics, University of Tasmania,
Box 252C, G.P.O., Hobart, Tasmania 7001, Australia
Abstract
Observations with high time resolution (<I ms) have been made in the frequency
range 30-82 MHz. During June 1979 several thousand bursts with bandwidth -100 kHzand duration at a single frequency of about 50 ms were recorded. Their frequencydrift rates were between -1.5 and -5 MHz s
-1. They belong to the classification
"fast-drift storm bursts- (FDS bursts). Approximately 1% of the bursts in thepresent observations showed bands or fringes in the frequency-time plane almostparallel to the time axis and with a frequency separation of about 100 kHz.Figures I and 2 show examples of FDS bursts and FDS bursts with fringes.
In this discussion the presence of fringes and their short dV iation ata single frequency are considered.
43
2-6-79
40-
I
)z
L 39-
02h07m35s UT 38s
Figure 1. Dynamic spectrum of FDS bursts. A fringed FDS
burst can be seen above 39.6 MHz a little after 02 071381.
I. FRINGES
We consider two possible generation mechanisms.
(i) The fringes may be due to the solar radiation having a strong linearlypolarized component, being subject to Faraday rotation in the corona andionosphere and being observed with a linearly polarized antenna. From thefringe spacing the path integral D = fNBiidk (where Ne is electron density,Bt is the longitudinal magnetic field and ; is the path length) was estimated.
It was found that near 40 MHz D = 3.8 *0.2- 10" C m-:. From simultaneousionosonde data the contribution to D from the ionosphere wasDI 4 .1012 G m
- 2, leaving the residual, presumably of solar origin, of
DC 3.4 ±0.2 x1013 G m- 2
. This corresponds rc a total Faraday rotation inthe corona of 503 ±30 rad at 40 MHz.
(ii) The second alternative considered involves an idea from a theory (Takakuraand Yousef, 1975) for the generation of Type IlIb bursts. According toMelrose (1974) the brightness temperature T of a source of fundamental plasmaemission is related to the electron density gradient as follows:
T - exp(;iL- 1) ,
< T >A
where C, f2
_p2 cm- ,Ti
44
40-
O2hlOmOgs UT 0g52Omsec
5.-
~L .I
02hlOmOs UT 11S
U-
Wiilll ! !I !il 1i! i
I ! I i l
02hO7m57s UT 56s
Figure 2. Dynamic spectra of solar radio bursts which show fztnges. The spectrumof each burst is shown tw-., recorded with a 256-channel filter bank (left) andwith a time expansion sweep frequency analyser (right). The sweep frequency analyserhas a time expansion ratio of 1500 and simulates a 1500-channel filter bank analyserwith individual filters of 2 kHz bandwidth. Each sweep of the analyser is 20 meapart and displays the power spectrum. All bursts were recorded on 1979 June 2.
45
where f is in megahertz, Ti is the ion temperature, Tp is the effectivetemperature of plasma waves confined in a limited range of solid angle A,and L is the effective distance over which amplification is possible. Now
L -grad Lp
where up is the plasma frequency and ,i - Ne . Thus the brightness tempera-
ture is strongly dependent upon the electron density gradient in the source.An electron stream travelling through a region of the corona with a spatial
periodicity in electron density might therefore generate fringed burstssimilar to those observed. Figure 3 shows the observed profile of brightnesstemperature with frequency for a pair of fringes and the profile expected
from a sinusoidal variation in electron density. The observed fringe separa-
tion implies a wavelength of the density variation of 1000 km, assuming the
variation is superimposed upon a 2 Baumbach Allen coronal density model.
Tb
38.5 FREQUENCY (MHz) 38.7
Figure 3. Variation of brightness temperature Tb(arbitrary, linear scale) with frequency for a pairof fringes. The solid line gives the observedvariation and the dotted curve is derived from asinusoidal electron density variation as described
in Section I(ii).
46
2. DU:RATION
Figure 4 shows the distribution of fringe durations. The burst duration sets
an upper limit to the size of the source of emission. A duration of 33 ms limitsthe source size to 104 km The short duration also sets a limit on the amount ofscattering of radiation within the corona. Numerical calculations of the expecteda:nount 01 scattering give minimum burst durations. Riddle (1974) has calculated
the expected time smearing of short bursts of 80 MHz radiation. His calculationsgive minimum durations of about 0.1 and 0.5 s for emission at second-harmonic andfundamental plasma frequency respectively. In view of the short durations measured
in the present observations, it seems important to repeat these calculations for0 MHz radiation.
30
20
E-oZ
10
20 40 60 80 100 120 140
Duration (msec)
Figure A. Distribution of durations of the fringes in FDS bursts.
It should be noted with regard to coronal scattering that the short duration
of FDS bursts suggests that the radiation must be emitted at the second harmonicof the plasma frequency. However, the coronal irregularity model proposed abovefor the generation of FDS fringes requires radiation to be at the fundamentalplasma frequency. In order to resolve this problem further study of these bursts
must aim to determine which harmonic is being observed.
47
References
Melrose, D.B.: 1974, "A relationship between the brightness temperatures for
Type III bursts", Solar Phys., 35, 441.
Riddle, A.C.: 1974, "On the observations of scattered radio emission from sourcesin the solar corona", Solar Phys., 35, 153.
Takakura, T., and Yousef, S.: 1975, "Type lIb radio bursts: 80 MHz sourceposition and theoretical model", Solar Phys., 40, 421.
Discussion
Stewart: The Faraday rotation interpretation assumes the stria bursts are linearlypolarized. We know that Type Illb stria bursts are unpolarized so I suspectthat the same is true for the fine-structure bursts you described.
McConnell: We have set up an antenna with crossed dipoles sensitive to twoperpendicular linear polarizations to try and resolve the question. We havenot vet observed any of these bursts in this way.
Nelson: Given that you have an estimate of the size of the source, have youcalculated its brightness temperature?
McConnel I: The flux from these bursts is comparable with that from the thermalsolar background. A source size of 10' km then implies a source temperature ofOI11 K.
Dryer: Your assumption of the sinusoidal density profile suggests that fast-modeM IID waves may have been propagating upward.
McConnell: Yes, the implied period of such a wave, with wavelength of "5000 km,would be 1-I0 s, depending on the model corona used.
Croft: Is it possible chat nothing is travelling but that the longer delay atlower frequencies is due to dispersion in the medium, and that the source was asingle, wide-band outburst?
,'lcConnell: We considered that possibility, but the delay at lower frequencies doesnot follow the correct law. Furthermore, we observe some bursts which show areverse-frequency drift.
48
5. The Mechanism of theU-Shaped Spectrum of Type IV
Solar Radio Bursts
Li Chun-sheng and Zheng %Wing-wuDepartment of Astronomy, Nanjing University
Nanjing, China
Yao Jin -ingPurple Mountain Ohservator%
Academia SinicaNanjing, China
Abstract
A mechanism is suggested to account for the U-shaped spectrum of the TypeIV solar radio burst during the explosive phase of a proton flare. The U-shapedspectrum consists of two parts, namely the part of Type IV-, burst and that ofrype IV-dm burst. Each of them is believed to be produced by gyro-synchrontronradiation of non-thermal electrons (100 keV * c < 2 Mev) being accelerated simul-taneously with protons in the explosive phase and gyro-resonance absorption ofthermal electrons.
Numerical calculations are presented for the flux densities of Type IV- ,hurst and Type IV-dm burst during the 1972 August 7 event, assuming a cylindricalvolume of these burst sources. It is shown that the computed results are con-sistent with that of observations. In addition, we also discuss the relationbetween the U-shaped spectrum of a Type IV radio burst and the occurrence of aproton event.
49
I . ... ..-
I. INTROI)UCTION
The peak flux density spectrum of Type IV solar radio bursts is generallycontinuous broad-band, particularly those associated with PCA (Polar Cap Absorp-tion) events having a distinctive U-shaped spectral signature (Fig. 1). Thertis high intensity in the meter range, a minimum at 500-4000 ',ffiz, a peak nearW1 GHz, and a decline toward higher frequencies. Such spectra were first con-structed by Castelli et al. (1967) who showed that this t-shaped spectrum is acharacteristic feature of all flares producing strong proton fluxes in space.Since, however, a well-developed Type IV burst is a typical characteristic ofall proton flares, this spectrum shape is in fact characteristic for all strongand well-developed Type IV bursts. A series of well-known -shaped spectralcriteria for proton events were formulated (Sarris, 1971; Cast,.lli and (:uidi,197o). These criteria, being used as a v s-no frec isting to.l 1 or prtin l ,nt,clearly had predictive or warning value, since the delay from the ohs-rvance ofthe burst (and the identification of the spectral signature) to the first onsetof the particles at the earth could be anywhere from 10 minutes to 20 hours.
2I, I
1 00,000-
E
10000 - AUG71521 UT
I000
300 1000 3000 10000 30000
FREQUENCY - MHz
Figure 1. The U-shaped spectrum of Type IV solar radio bursts on 7 August 1972.
During the period 2-7 August 1972, four unusually intense solar flare-burstevents occurred in the McMath plage region 11976. All of them were associatedwith the U-shaped spectrum and proton events. Moreover, those on August 4 and 7were ascertained as ground level events (LE's). During this period extensiveobservations of the spectacular solar events were made by all available meanssuch as spacecraft, balloons, lunar-based instruments, and rockets in additionto numerous ground-based solar observations both in optical and radio regions.
The data of radio observations over wavelengths ranging from millimeter tometer for the U-shaped spectrum of the Type IV radio bursts on 7 August 1972 attile same place (Sagamore Hill) and the information of related observations wereprovided (Castelli et al., 1973).
50
We also utilize the observational data published in Solar Geopihysical Data7 August 1972 by Academia Sinica. It is, therefore, appropriate that we utilizethis extensive data set to study the physical mechanisms of the U-shaped spectrumof Type IV radio bursts.
2. MODEL OF SOLAR FLARE-BURST SOURCES PRODUCING TIlE U.SII.PED SPECTRUM OFTYPE IV RADIO BURSTS
In this study, the U-shaped spectrum is considered to be formed at the peakradiation of Type IV radio bursts during the explosive phase oi a proton flare.The range from the valley on the curve of the U-shaped spectrum to higherfrequencies is the spectrum of Type IV-. bursts. The range from the valley tolower frequencies is that of Type IV-dm bursts. These two spectra are believedto be produced by the sources of Type IV-1- bursts and Type ]V-dm bursts,respectively.
During the explosive phase of a proton flare, the ejected plasma cloud,frozen in the magnetic field, seems to be the source of Type lV-dm bursts. Themagnetoactive plasma cloud captures numerous high energy electrons ejected upwardcontinuously with arbitrary angles from the solar flare explosive region.
In the course of gyrating around magnetic field lines, the high energyelectrons emit electromagnetic waves, generally known as gyrosynchrotron radiationto form the Type IV-dm bursts. Similarly, at ,i:. ,ame time those ejected downwardcontinuously will also send out gyrosynchrotron mission to form the type IV- .bursts.
All the high energy electrons mentioned above ant those higih energy protonsescaping into the space during the explosive phase of a protou flare may beaccelerated simultaneously by the same accelerating mechanism, ihus a closecorrelation may exist between the U-shaped spectrum of type IV radio bursts and
the proton events.
The model described above is analogous to that of Sturrock (1968). Forsimplicity, the magnetic field due to a bipolar sunspot is illustrated by twoindependent axially symmetric dipole fields which have opposite polarities plottedin Fig. 2. The magnetic field-strep, th along the axis with height is given by"akakura and Sca lise (1970),
d4-h
where H is the magnetic field strength at the photosphere and is set to 3000 G
(Castelli et al., 1967; Solar Geophysical Data, 1972). H is the magnetic fieldstrength at the height h above the photosphere. The depth of the vertical di-pI I e Iw Ihe pIhotspher,., d, equa s 1.5 x 104 km in the present model(Tikikura alnd S-,i Iiste, IliO).
51
120* I Vd., SOURCE
FLARE EIIPLOSION -REGION
Figure 2. Model with a bipolar sunspot for solar flare-burst sources.
In the sources of Type IV-lj bursts and Type IV-din bursts, the magnetic fieldlines are approximately antiparallel to each other, but all are taken to be parail nto tile photospheric surface.
For computational simplification, the volume of these bursts' iources areassumed to be cylinders, and their bases are all parallel to the V otosphericsurface. Thus for the active region McMath 11976 (N14', W36
°) on i August 1972,
the 3B flare meridian distance can be adopted as 30'. If - represents the anglebetween the direction of observation and the magnetic field in thie source, thenit will be bO° for th, type IV-. hurst ourrce and 120' for Vpt lV.-dm burst -olirtin this case. Hence it follows that the quasi-longitudinal propagation of radiowave from the above radio sources under consideration is valid.
3. TIll- COIPI'TIN; FORNI I, %F" FOR I -SlI..PED SPECTRUM
According to the theory of radiative transfer, the spectrum of the emissionfrom the solar radio source is given hv
ds(f,,,) = r(f, ) e- (2)
where f is the frequency ot emission under consideration and ds(t,-) is the fluxdensitv received at the distance R to the earth due to the gvrosvnehrotron emissionwith volume emissivity n (f,u) from a volume dv. 1(f.0) is the optical depth of thebackground plasma in tile source for a volume dV and is supposed to be due mainlyto thermal gyro-absorption in the present model. (The self-absorption in the
52
sources is assumed to be small).
3.1 The Volume Ejnissivits
We will use the gyrosynchrotron volume emissivities in a magnetoactive plasmagiven by Ramaty (1969) and the corrected ones by Trulsen and Fejer (1970) and Koet al. (1973). We consider the distribution of pitch angle of energetic electronsto be isotropic, the energy distribution N(E) Gc
- !' dc, and the non-thermal
electron distribution in the burst sources to be homogeneous. Then we obtain thevolume emissivities 1,2(f,6) of the two polarization modes (extraordinary and
ordinary waves) for 0 1 90' as follows:
~2f(f',) = GTe f "1,2 s ' (3)
1,2( ' ) 2, 1,2~ nI
2 2 2"1, a 2h co 0 2 l+p -1) d
a-b 'o 61 P 1,2 1+ 2 d
a p sin J .' (x) [2r,.s 'I +p-
s2(1 +p2 sin2., )_2
1 2
s COS " n
b= -s ,1 +p 2 sin
2 J ln)
s cos -sin
f fs fH (f in M-iz and H in Gauss) (4)
-2f(f2 -f ) cos e
p(5,1,2 f 2fH sin
2 U± [f 4fH2 sin 4 +4f 2(f
2 f 22 cos2 1/2
p
f (Mhlz) = 9 x io ', (6)
wherv p = ---- (mo is the rest mass of an electron; my is the mass of an electronmoC0
moving with velocity v; and c is the velocity of light); F is the spectral index
of the energy distribution; NT (in cm- 3
) is the number density of thermal
electrons. The suffix 1 and 2 or the upper and lower signs in equation (5)
correspond to the extraordinary and ordinary waves respectively; .n (x) is the
Bessel function of integral order n, the prime on the Bessel function denotesdifferentiation with respect to its argument, C is the proportionality coefficient,
fH is the gyro-frequency of an electron in a magnetic field H,
53
f is the plasma frequency and 1,I, 2 is called the polarization coefficients of
the two characteristic waves for e = 60'. If f > f and f f H' one obtainsp
,,2 = Fl, in which O!71 = -1 and t02 = +1 will correspond to right-handed and
left-handed circularly polarized waves, respectively.
3.2 The Thermal Gyro-Absorption
The absorption coefficient of the thermal gyro-absorption has been given bvseveral workers. Hert, the formulae reduced by Zlotnik (1968) are used since these
seem to be convenient for the computation.
For quasi-longitudinal propagation (for the above-mentioned cases - 60' and120'), if f i and s - 2, one can obtain from Zlotnik (1968) the opticalthickness of thV gyro-resonance absorption layer in the plasma with inhomogeneousmagnetic field as follows:
2s f2
= 2 2s fp .2(s-1) 1)2 ,.)2(s-l) (7)- ___ 1(1± C s 7s 2's! f T If csJ)(i
where '= -K- T is the electron temperature, K is Boitzmann constant, andr moc '
IH is the characteristic length reflecting the significant change of magnetic fieldalong the line of sight. The upper and lower signs correspond to the extraordinarywave (I =1) and ordinary wave (j =2) respectively.
According to those mentioned above and, together with the configuration of themagnetic field in Fig. 2, the gyro-resonance absorption laver for individualharmonics is parallel to the photosphere. If s(f/f H ) represents the harmonic
number, tien after substituting s into equation (1), we can get the height habove the photospherV for the given frequency f and harmonic number s, i.e.,
2.8 110 1/3
1 = d 8 I - 1 .(8
l-tirtLhrmorc, we can also get the thickness of the emitting laver
d 2.8 11o 1/3 -2/3 (9).As = 3 f
3.3 The Stectrum of Gyro-Sychrotron Enision From the Assumed Cylindrical Radio Sources
The spectrum and polarization of the emission from the assumed cylindricalradio source (Fig. 3) can be computed from Equations (2), (3), (4), (5), (6), (7)
and (9) for two modes of the wave.
54
TO THEEARTH/
CMP
CENTRAL / 30* DIRECTION OF
MERIDIAN DISTANCE MAGNETIC FIELD
9 * 60"
nS-5 'IS2T25
Ah4l { ')14 7)24 5.• 4 TI4 , T24A h'7{ 1 3 "7 2 3 S T 3 I 2 3
Ah2 ~'12 '722S 2 r, r2So2 TI2 ,1"22LAhl {11 'r/
H r_________________ S_ 5. I 11 ,tI 2
l ---- -H2
w PHOTOSPHERE E
Figure i. The distribution ot gyro-resonance absorption laver and emitting laveri i ,i vlinudrical m,,d,[ I l'VpL I%'-. radio hurst-,sour v..
For the givenl frequencv f, if "Is and "2s represent respectively tile volume
emissivity of extraordinary and ordinary waves between thii gvro-resonane absorptionlaver of the s-th harmonic and that of the (s +l)-th harmonic, and if is and -2s
denote the optical thickness of the gyro-resonance laver of the s-th harmonic of the
extraordinary and ocdinary wave correspondingly, then under tile condition considered
in tile present model (see Table 1)
i (), when s < 6 .
Hence the absorption factor exp (-Is) is necessary here only at s smaller than 5.
hus, for a 5, tile flux density received at the earth due to the emission fromtile corresponding volume is tile summation of the contributions from two mod-. atwave in each emitting la.er below tile s =5 absorption layer according to the
radiation transfer equation (2). However, the flux density due to the volume
emissivity contained from the s 5 absorption laver to the upper boundary, s =SP
of tile radio source is determined by the corresponding integration from thef
boundary limits: s =5 to s =.u = ----- I in which H represents the magneticlUP 2.8 ifu up
field strength at tile Lipper boundary (Fig. 3).
Then the total flux density S(f) due to the entire radio burst-source for
' =6)' becomes
5
S(f) = s1(f) + S2 (f) (10)
where SI(f) and S2 (f) represent the total flux density contributed by the extra-
ordinary and ordinary waves, respectively, with
222
S f e YO f 1 e(113+ C14 + E1 5)cR
2 2 212e
- T + +1 Ah d-1
Ah 313e' 14 + 115) + Ah 4 , 1 4 d- 5
+ d (2.8 o)u1 /3
pS5/3
ds3f/ 1/3 f~ ' *15 (1
5
222G 7 e 7 o f + 2 3 + 2 +
2
S2 (f) = cR 2 21 22 24 25
+ I "h 2, 22, (113 + t24 + 25) + 1 tAh 3: (,24 + 125)
+ 7 CR21 ds,
424 (12)5
where r is the linear radius of the radio burst-source observed,0
S13 = SI(s, e = 60°) and (13)
S,4 - S2 (s, 6 = 600)
in equation (4).
In addition, one can obtain the number of nonthermal electrons per unit volumein the radio burst-source
tm ax -r
N = G J di , (14)
min
hence the total number of non-thermal electrons is NV, and . and m- ismn max
determined by p1 and p2' respectively.
From equation (5), ine can obtain for H - 60': 1,2 = fl in which
= -1 corresponds to the extraordinary wave reduced to a right-handed circularly
56
polarized wave, thereby the flux density may be repres--nted by S1 = SR . The quanti-
y .= +1 corresponds to the ordinary wave reduced to a left-handed circularly
polarized wave; the flux density may be represented by S2 = S , hence the polari-zation degree for Type IV-. radio burst is given b 2
SlM - S 2 (f)
S (f) + S2 (f) 15)
For = 1200, a01 2 =
1, i.e., the sense of circular polarization is reversedfrom o = 60', thus' the polarization degree of Type IV-dm radio burst is given by
S 2(f) - S (f)
S 1l(f) + S2(f ) 1
4. CONCLUSIONS AND ANALYSIS
As stated above, we have cited some physical quantities (see Table 1)published by Castelli et al. (1973), Sturrock (1968) and Castelli et al. (1967),to compute the U-shaped spectrum of Type IV radio bursts at 1521 UT 7 August 1972.
The computed results are listed in Table 2 and the observation date of U-shaped spectrum is plotted in Fig. 4.
Table 1. Physical Parameters of Type 1V Radio Burst-Source and Type IV-dm RadioBurst-Source
Type IV Type IV-dmPhvsical Parameter Radio Burst Radio Burst
Photospheric field magnetic 3 3Ho (Gauss) 3 x 10 3 x 10
Angular diameter 1.6 3'.0Height of the lower boundary 4
h (km) 1.5 x 10 1.13 x 105
Height of the upper boundaryhu (km) 6 x 10
4 1.7 x 10
5
Temperature in the plasma of 6 6source 5 x 10 6 x 10
Number density of thermal electrons 10 8(cm
- 3) 10 2 x 10
Characteristic length 10l.H1(cm) 4 x 10
9 4 x 10
0
Angle - (degree) 600 1200Minimum momentum of non-thermal
electrons P, 0.66 0.66Maximum momentum of non-thermal
electrons P2 4.70 4.70Index of energy spectrum F i.0(*) 1.0(*)
(*) The value of I is adopted from Castelli et al. (1967).
57
Table 2. Parameters Related to the Non-Thermal Electrons in tile Radio Bursi-Source
Type M; Radio Type IV-dm RadioParameters Burst-Source Burst-Source
Number density of non- 3 3 3 3
thermal electrons 1.0 x 10 (cm- ) 1.7 x 10 (cm
Total number of non-thermal electrons in 32the radio source 2 x 10 1.4 x 10
4"1 0 105
x
X
- 4- 10
z
x
-J
w
102 O3
OFREQUENCY (MHz)
Figure 4. The comparison between the computed U-shaped spectrum (solid line)and observation (dots) depicts the valley between the Type 1%'-. andTp\'p IV-dm radio sources.
Although we simply assume the volumes of Type IV-. and Typ, lV-dm radi,,olr's t" b" vI indrs in the models mentioned above, we found, nevertheles.,, that
computed U-shaped spectrum of 1%'pI lV radio bursts is consistent with thatobserved. Thus, we can draw the following conclusions:
1. During the explosive phase (the second acceleration phase) of a protonflare, the effective emission mechanism generating the U-shaped spectrum of TypeIV radio bursts is due to gyrosynchrotron radiation and the effective absorptionmechanism is due to gyroresonance absorption, so that the data concerned in thecomputed model may approach Che real physical param'eters in the radio burst-sources for this event.
2. The peak frequencies of Type IV-i, and IV-dm radio bursts are determinedby the equation f = (3-4) x 2.8 H. , in which HI represents the magnetic field
mas
of the lower boundary of the radio sources. The valley between these two peakfrequencies appears near (2000-3000) M,2 which is the frequency range betweenthe lower frequency part of Type IV-, radio burst and the higher frequency partof Type [V-dm radio burst. The corresponding emissions in this frequency rangeare generated In the region between the upper boundary ot Type [V-dm burst-source
58
and I h lower boundar'.o I1 .p, IV'-. burst-source, i. e. ,the flare explosive regionor the annihilhat ion region of the magnetic field during the explosive phase of theproton flare. InI this region the magnetic field strength is so weak that it mustlead t, the conclusion that the volume emissivity of gvrosyncitrotron radiation
Mulst decrease rapidly. Moreover, the lower frequencies' emissions of Type IV-.bursts are not only afftected ily the gyroresonance absorption hut also suppressedhv tile background plasma itself (see Svestka, 1975, for a discussion of theKazin effectL), whereas the higher frequencies emissions of Type, IV-dm hurst arerap idly decreased due to the fact that their emit ting volumes diminished rapi il>withi increas ing fre0qucIies1. As a consequence oi these combined effects,* a valleyis termed near tihe f reqjuency range 20)00- 30(00 Mliz on the peak flux dellsfty spect rum
ofT type 1% radio bursts. ltccasionallv, a peak emissionl Ima'. h observed at the
sautle I retiencv range dllring tihe proce-ss of somei t iirc-bllrst . 111L "it 11 tIll' I
between these f lor-bursts andi thiose produicing tite V'-shaped spec trutm ;,&'% heexplained ite tite eject ion of conlsiderable amnount of material. irobable , Lteformer itas neititer developed vet to tile explos ive pihase nor ejectcd ;uffi,_it-ntmioter ial to form tite sources of 1%1,,- IV- and 1.1pe IV-dm bhursts ISe a
3. fihe foregoing di SCLI~si n ii u!.t rates the suggestion t hat t ih U
spect rumn of tiepu 1V rod in bur-sts, const itties radio ev idenct olf t~i, rot en !L., rdevelopment to the second accele rat ion pitase. turing this 'hoe t11C II-etiron-and protons are accelerated imtutitanuous '. to itigh energi 1. artijetlar I..
ICC t rills ma'. be accelerated to itigh energ.' ( >eV) , hut th~ e can heI iCCh lt-
only to 10-100 KeV during tite f lasit phtase (tite first acceIleration 'ba-) Z1-discusse-d be Stutrrock (1974).
it is shiown Me sp~ace Oi., -reion, thtat tile second accelera iott ;ha-,~ n1 1-
aitout -1t3 minlutes after the t'JI.1 sh pase, and the energy9 of p'roteusop )lK'
100 MIAV even I BLA * as noted Ibe Caste 11i and ;u id t, 19 7t) .ito,, it Lcan 1,,0C11 ThaOt t ite oec ir rences oi th Ito -shaped spect rum 'lust clear'. citrr l at, witli
these Of the p r-oton eIVIelt 5.
ca.stelli and 61ii- I l',7,) pointed out that tile i'-itapod sl'ectrum~ crit, rio
used I or pred ict ion Lf1" ;Toten events) came to he identified mainly withi printcipal
'C';i.e. , thtose with iestired r iometer absorpt ion h2db at 30 1Hz. However, tite
sat eliito detectois cal de tect nmerouls mucih weaker proton events produced by, small
1 ~ ar. lies sub;, 4i I i tf ten produice nei titer great ejict ion of material nor
I -hiaped pcect r1'm 1ee t, 1 ,*1 thi- L -staped spectrum is associated witht a proton
e-vent , Ih, 'ro5.iI 'it T th hut wen tito proton evettS (including tite proton
event' Jet >1 tid 1% ti i te ) oCr witht tue U-sitla:ted s poetrIu appearing
bet10-1 ore ,ltd, t.j H i, ttlahi.11t', is ''re low. The albo'.e descript ion mav account f-or
t Ile r, Lit ott i't,ee tin- protonl cvtts and tite I-shaped spectrum studied I)% Moarris
I17)Withi satellite- iCxplor, r 3 i and Vt.)
iron (Oslo,1) -*: al. (19731, we Can sue that U-shape'd Spectra associated witih.proton ,,*cttt or grouind level e'vent were ob-served, w.itihout except ion, on 2, 4 and
7 August 1972. These e~vents ttbv iotisl\ indicate t iaot t he electrons and protoensaCeCIerat ed s imttl taleous l% to itight energ ies arc vore large in numtber during thle
,xpies ivc pitasc of proton flares, stt titat tile itigit energy protons escaping to
nt,,rplanc'tar. spa.ce, even1 to cir1tit, !beC0::, larger itt tttm'L,'r.
For the event on 7 Allitlst 1972, ti,, pee trun of 1%pe i\'-. rai, 'ittrst
sitowed t tat tilt mic "owaec part cot harder attd itarde-r (i.e., tie spect rum curveget flatter and f1;lttr) witht tile lapskc of time.l At 1 -18 1 '1, tite pa rt of the
spectruim icieir Thatn the peak ft iluene'. arriee'i at it, hardest Ilev (Fig. Sforte mintits lot, r, .1 t C Cirrkti t 1000)1 ['1 is Cast.,li i t al., 1973).
59)
40000
-20000-
z- 1522
-4 151800 0 -
1,
/ "<Z 15155L.J
X 4000-
2000
1000 I I I2000 4000 10000 20000 40000
FREQUENCY - MHz
Figure 5. Flux density spectra at different times during the event of 7 August1972. Note that spectrum above f max is hardest at 1518 UT.
This phenomenon is believed to be due to the fact that the electrons and
protons are accelerated to the high energies; while the gyrosynchrotron emissionscaused by the high energy electrons are strengthened, the high energy protons
escaping to the interplanetary space are increased. Therefore the evolution of the
microwave part of the U-shaped spectrum of Tpt' IV radio bursts can be used for
proton event forecasting.
4. Finally, the discussion in §3.4 shows that the radio emissions fromTlype IV-, and FTpu, IV-dm sources are all circularly polarized waves under quasi-
longitudinal propagation. The polarization degree for each radio source computedIrom equation (15) or (16) with the data listed in Table I are shown in Table 3.
One can see trom l'ahle 3 that the sense of polarization reverses at the frequency
21)00 or 30001 M z . \t the higher frequencies the right-handed circular
polarization is dominant, but at the lower frequencies the left-handed
one is dominant. This phenomenon arises from the fact that the magnetic field
lines in the Type IV-, burst-source are anti-parallel to that in the Type IVbIur[ s-seul rc e.
Since we adopt a simple cylinder as the volume of the burst-source in our
computat ion, the computed results are not accurate enough for precise comparison
witih observations and should be used essentially to gain insight for Lhe overall
problem.
60
Table 3. Variation of the Polarization Degree with Frequency for Type IV- .and Type IV-dm Radio Sources at 1521 UT 7 August 1972.
Type 1V Burst-Source Type IV-dm Burst-Source= 600) (= 120o)
f(MHz) i(f)% f(MHz) ,(f)%
2000 -28.6 300 -20.33000 +27.6 400 -25.95000 +31.3 500 -30.9
10000 +31.7 600 -33.413000 +33.2 700 -31.815000 +32.2 800 -30.520000 +29.7 900 -20.425000 +28.0 1000 -28.530000 +26.7 2000 -24.035000 +25.6
Acknowledgments
We would like to express our appreciation for discussions with ourcollagueCs Zhou Ai-hua and J]iang Shu-ving.
References
Castelli, J. P., Aarons, J., and Michael, G. A.: 1967, "Flux densitymeasurements of radio bursts of proton-producing flares and nonprotonflares", 7. Geophys. Res., 72, 5491.
Castelii, J. P., Barron, W. R., and Badillo, V. L.: 1973, "Highlights ofradio activity during disk passage of HcMath plage no. 11976", inHelen E. Coffey (ed.), Collected Data Reports on August 1972 Solar-
Terrestrial Events, Part I, (UAG-28), U. S. Department of Commerce,NOAA, Boulder, p. 183.
Castelli, J. P. and Guidice, D. A.: 1974, "Solar bursts at - 2 cm onJuly 31, 1972", 7. Geophys. Res., 79, 889.
Castelli, J. P., Carrigan, A. L., and Ko, H. C.: 1974, "Spectral associa-tion of the 7 August 1972 solar radio burst with particle accelera-tlon", in G. Newkirk, Jr. (ed.), coronal Disturbances, IAU Symp. 57,D. Reidel, Dordrecht, p. 143.
Castelli, J. P. and Guidtce, D. A.: 1976, "Impact of current solar radiopatrol observations", Vistas in Astronomy, 19, part 14, 335.
61
References
Ko, H. C. and Chuang, C. W. : 1973, "The relation betweei the radiatedand received powers due to gyrating electrons in ambient plasma",
Astrophys. Letters, 15, 125.
Ramaty, Reuven: 1969, "Gyrosynchrotron emission and absorption in amagnetoactive plasma", Astrophys. 7., 158, 753.
Sarris, E. T.: 1971, "Study of solar flares by the satellites Explorers 33and 35", (AD-731686), MS thesis, University of Iowa.
Solar Geophysical Data: August 1972, edited by Academic Sinica.
Sturrock, P. A.: 1968, "A model of solar flares", in K. 0. Kiepenheuer(ed.), structure and Development of Solar Active Regions, IAU Symp.35, D. Reidel, Dordrecht, p. 471.
Sturrock, P. A.: 1974, "Particle acceleration in solar flares", inG. Newkirk, Jr. (ed.), Coronal Disturbances, IAU Symp. 57,D. Rediel, Dordrecht, p. 437.
Svestka, Z.: 1976, Solar Flares, D. Reidel, Dordrecht.
Takakura, T. and Scalise, E.: 1970, "Gyrosynchrotron emission in a mag-netic dipole field for the application to the center-to-limb varia-tion of microwave impulsive bursts", Solar Phys., 11, 434.
Trulsen, J. and Fejer, J. A.: 1970, "Radiation from a charged particle ina magnetoplasma", J. Plasma Phys., 4, 825.
Zheleznyakov, V. V.: 1970, Radio Emission of the Sun and Planets, Pergamon
Press, Oxford.
Zlotnlk, E. Ya.: 1968, "Theory of the slowly changing component of solar
radio emission. I.", Soviet Astron., 12, 245.
Discussion
Smart:. Your model for generating the U-shaped spectral signature is very
Interesting and has the potential of satisfying the theorists' principalobjections (i.e., that two different emission levels in the solar atmosphereare required). However, my question is directed to Professor Wu. Does your
hydrodynamic model of a shock propagating through the helmet strearer con-
figuration generate a pulse (increase) in magnetic field that would act asa reflecting barrier for electrons in the top of the streamer such that therewould be an emission level there?
Wu: Yes, I believe so. For example, we show the density enhancement rightbehind the shock in which the emission is enhanced. Magnetic flux enhancementswould also be present. Conditions appropriate for either (or both) plasmaoscillations and gyro-emission are predicted by our model.
62
Discussion
Giovanelli: On the Sun, an unstable condition leading to an explosion cannotexist with such a simple model (and of course it requires a finite conductivity).Hence in reality the field structure - and the problem to be solved - becomesmuch more complicated.
Dryer: I would like to comment that Dr. T. Yeh is currently looking into the
reconnection model associated with such a complex magnetic topology such as thatproposed years ago by Dr. Giovanelli. So'I would agree that the simple sketcheswhich usually appear in the literature are probably too simplistic and possiblytopologically untenable.
63
0!
6. The Proton Flare of 1976 April 30
S. 1. Avdjushin and N. K. PerejaslovaInstitute of Applied Geophysics,
Hydrometeorological Service, Moscow, U.S.S.R.
V. V. Fomichev and 1. M. ChertokSolar Radio Laboratory, IZMIRAN, Moscow, U.S.S.R.
Abstract
It Is possible to use solar radio data to predict the flux of protons froncertain flares. The flare of 1976 April 30 gave us an opportunity to compare suchpredictions with the satellite measurenents of particle fluxes.
I. INTRODUCTION
Rapid development of the active region Mcrlath 14179, which began on April 29,.culninated in a proton flare at 20h47 1 UT on 1976 April 30 of importance 1B at 0S, 46 W (Solar Geophysical Data, 1976). This flare was a source of protons with a
65
wide range of energies (Solar Geophysical Data, 1976; Avdjushin et al., 1977). Itwas followed by an intense radio burst at centimetre, decimetre and metre wave-lengths. In such events it is possible to use the radio data to predict the flux ofsubsequent protons. This particular event gave us an opportunity to compare suchpredictions with satellite measurements of particle fluxes.
rhe frequency spectrum of the maximum flux densities of the radio emission(Figure 1) was found from the data at fixed frequencies reported in SolarGeophysical Data (1976). The fluxes exceeded 1000 s.f.u.* at metre and centimetrewavelengths and had the U-shaped peak power spectral form typical of proton flares(Castelli et al., 1967).
2. METHOD
The maximum flux density of microwaves at a frequency of 3 Ghz may be used toestimate the expected flux of accelerated particles. According to Akinyan et al.(1977) the maximum intensity (J) of proton flux with energy E > 10, 30 and 60 'leVrecorded near the Earth and the radio flux (S) of the burst for longitudes between20 and 80 W are empirically connected in the following way:
J= J(S)- (1)
where J(S) is found from:
for E 10 MeV log i = 0.0295 S0 .
5 8
for E > 30 'MeV log J - 0.055 S0.5 -1.O
for F > 60 '1eV log J = 0.243 S0-338 -2.0 if S 4 3000
log j = 0.1 S0 " 4 5
-2.0 if S > 3000
(S is ,iven in s.f.u.) and t is a parameter computed from the intensity ofmetre wave emission which accounts for the ease with which the accelerated particleqmay escape from the flare region.
In the present case the parameter S = 2000 s.f.u. (Figure 1). The intensityof the metre wavelength radiation at frequeacies ( 245 MhZ seems not to have exce-eded the value 1000-2000 s.f.u. Observations with the Culgoora spectrograph showthe radio emission at metre wavelengths included bursts of Type II and IV withintensity of importance 2. According to Akinyan et al. (1977) such an intensityis not enough to facilitate the escape of particles from the flare region. Inthi- case they give - 0.7. Using all the data above and euation (1), te obtainthe following values of maximum proton flux in particles cm
- s - 1
sr-
I nearthe earth: J - 170 for energies > 10 MeV, J - 21 for energies > 30 'IeV and J 9for energies > 60 MeV.
To construct time profiles of proton flux one also needs to know the delays be-tween the maximum of the microwave burst, and both the beginning of Increase, Atj,and the maximum, At 2 , of particle flux, and a parameter, T, describing the
*1 s.f.u. - 10- 2 2W m- 2
Hz- !
66
IdN
-J 10 3
LUCi0. S
02
102 103 104 105
FREQUENCY (MHz)
Figure 1. The spectrum of maximum flux densities of radio emissionaccording to the data reported in Solar Geophysical Data (1976).Note the U-shaped peak power spectrum typical of proton flares.
exponential decay of the proton flux. A statistical analysis (Akinyan and Chertok,1977) has shown that for heliographic longitudes 20*- 80 W the mean values of theseparameters are: AtI - 0.5 h for protons of all energies under consideration andt2 - 4 h for E > 10 MeV and 3 h for E > 30 and > 60 HeV. However, for phenomena
with relatively unfavorable conditions of escape, when the intensity of metre-waveradiation is <5000 s.f.u. and S ) 1000 s.f.u. the mean delays increase AtI up toapproximately 1.9 h and At2 up to approximately 9 h. The spread of these valueshowever, is high and the values may lie anywhere in the intervals 0.4 'At1 < 4.3 h,and 3.3 ¢At2 ( 13.4 h.
The decay of proton intensity with time follows a law which is approximatelyexponential, with mean time constant T = 9, 8 and 7.5h for E > 10, 30 and 60 1eVrespectively.
3. RESULTS
The proton fluxes following the flare of 1976 April 30 were recorded in thepolar region of the Earth's magnetosphere by the Soviet Meteor satellite (Avdjushin,et al., 1977). The intensity of proton fluxes with energies E > 90, 40, 25, 15and 5 NeV were measured.
In Figure 2 filled circles show observed maximum proton fluxes in differentenergy ranges. Open circles show the predicted proton flux for E > 10, 30 and60 IteV. Figure 3 compares the observed time profiles of protcas (points) with the
67
predicted time profiles (lines). The agreement between the observed and the predic-ted time profiles is generally satisfactory. In particular, the most importantparameter - the maximum proton flux in each energy range - differs from the obser-ved value by less than a factor of 1.5. The predicted decay rate is also closeto that observed. The time delays Atl (0.5 h) and At 2 (2.5 to 3 h), appear
to he lower than the predictions based on the comparatively weak metre wavelengthenission. This reflects the real condition of escape from the corona and theinterplanetary medium and the high spread already mentioned of Atl and At2typical in this type of event.
4. CONCLUSION
In conclusion, the parameters of the proton flux reaching the Earth after the
flare of 1976 April 30 could be satisfactorily estimated based on the informationcontained in the radio bursts at metre and centimetre wavelengths.
10 3 , , ,
St2
E
MV 1020
C1
00
100 I
5 10 15 30 50 100
ENERGY (MeV)
Figure 2. Protor fluxes following the proton flare 1976 April 30
recorded in the polar zone of the magnetosphere of the Earth byMeteor (Avdjushin et al., 1977). Filled circles are maximumintensities of proton fluxes with energies E >90, 40, 25, 15and 2 MleV. Open cicles are intensities of proton fluxes with
energies E > 10, 30, and 60 MeV predicted from radio observations.
68
10
1 - E > 5 MeV2 -E >15 MeV3 - E > 25 MeV4 -E > 40Me V5 - E > 90 MeV
10 2
E > 10 tMeV
3.
C0 1
Of 4
30 0
April1 May 1976
Figure 3. Time profile-; of protons: point,; re~resentexperimontal ,rofiles of proton data obs-erved by the !ieteor5atollite; solid lineq are calculated profiles based on radiola ta.
Acknowledgment
hils work forms part of the program of the Australian-Russian Science A 'ree-men t. Tlic author,; thank P.A. rhunran for assistance with the Enplish preseontat ion.
69
ReferencesAdin'yan, S. T. and Chertok, 1. M.: 1977, "Determination of the paratoetcrs ot
solar protons in the vicinity of the Earth from radio bursts. Ii. 're .pOr,
cefervnce functions.", Geomagn. Aeron. (U.S.S.R.), 17, 596. (English Transla-tion, Geonagn. Aeron. (U.S.A.), 17, 407.)
Akin'yan, S. T., Formichev, V. V., and Chertok, 1. M.: 1977, "Determination of theparaneters of solar protons in the vicinity of the Earth from radio bursts.Nt. Longitudinal attenuation function.", Geomagn. Aeron. (U.S.S.R.), 17, 177.(English Translation, Geomagn. Aeron. (U.S.A.), 17, 123.)
Avdjushin, S. 1., Pere jaslova, N. K. , Kulagin, Yu. I., Nazarova, 1. N., andPetrenko, 1. E.: 1977, "Observations of solar cosmic rays by ' teior' satvllit,,in March - 'Hay 1976", in 11. E. Coffey and J. A. McKinnon (eds.), CollectedData Reports for STIP Interval II 20 March - 5 May 1976, (UAG-61), U. S. I)e-partirent of Comnerce, NOAA, Boulder, p. 157.
Caste~li, J. P., Aarons, J., and Itichael, G. A.: 1967, "Flux density reasurementsof proton-producing flares and nonproton flares", J. Geophys. Res., 72, 5491.
Solar Geophysical Data: 1976, U. S. Department of Commerce, Boulder, Colorado, No.386.
70
70
!L , j .. .
7. Relativistic Solar Particle EventsDuring STIP Intervals II and IV
M. A. Shea and D. F. SmartAir Force Geophysics Laboratory
Hanscom AFBBedford, Massachusetts 01731, U.S.A.
Abstract
Using spaceship 'Earth' as a detector located at 1 AU, the relativistic solarcosmic ray events of 30 April 1976 and 22 November 1977 are compared to deduce therelativistic solar particle flux anisotropy and pitch angle characteristics in theinterplanetary medium. These two ground level events occurred during STIP Intervali1 and IV respectively - periods of time of coordinated and cooperative scientificefforts.
I. INTRODUCTION
The concept of STIP Intervals was adopted in 1975 by the membership of theSCOSTEP Study of Travelling Interplanetary Phenomena for the purpose of concentra-ted study of a specific time interval and the solar-physics phenomena that may occur
71
STIP INTERVALS
02.; ."o-
WN MUsJnl , 00
•. ..... .. .....
0~
Figure 1. Smoothed sunspot number from 1976 through 1979.The STIP Intervals are indicated at the top of the figure.The occurence of ground-level events is indicated by arrowsextending from the top of the figure, and the length ofthe arrows indicates the magnitude of the ground-level event.
during the selected period. 'lost intervals were selected in advance with primeconsideration given to radial or magnetic connection of various spacecraft. Eachinterval was designated for a two month period In an effort to maximize studies oftemporal and spatial phenomena without overly burdening the satellite tracking facil-ities.
Six STIP Intervals were designated between January 1976 and December 1979, asshown in Figure 1. This four-year time span included solar minimum in 1976 and therising portion of the 21st solar cycle as indicated by the smoothed sunspot numberalso illustrated in Figure 1. There were six ground-level relativistic solarcosmic ray events during this time period, with the events producing increases onsea-level high latitude neutron monitors ranging from two percent to 214 percent.Two of these ground level events occured during STIP Intervals. One was on 30April 1976 near solar mininun during STIP Interval 11, and the other was on22 November 1977 during STIP Interval IV. In this paper we will use the neutronmonitor measurements on spaceship "Earth" to compare these two relativistic solarparticle events and deduce the relativistic solar particle flux anisotropy andpitch angle characteristics in the interplanetary medium.
2. SOLAR PARTICLE PROPAGATION
Figure 2 is an artist's concept of a solar flare that is associated with theacceleration and release of high energy (E ± I MieV) solar protons from the suninto the interplanetary medium. Many of these particles travel along the interpla-netary magnetic field lines to I AU (i.e. Earth's orbit) and beyond, with thenaximun in particle flux usually measured along the interplanetary magnetic fieldline originating at the flare site. Since the large scale topology of the interpla-notary magnetic field is determined by "frozen in" tagnetic fields being transportedout of the rotating sun by the solar wind, the earth is connected to the sun by an
72
Figure 2. Conceptual illustration of the characteristics ofan anisotropic solar proton flux propagating along the inter-planetary magnetic field lines from the sun to the earth.
SUN -TEARTH LINEH
Figure 3. Conceptual illustration of an idealized "most favor-able propagation path" for fast and slow solar wind speeds.
30ARL1976 7LO
EARTH4(IMPS)
Figure 4. Position of Helios space probes relative to tht, sun-earthline on 30 April 197b. The flare occured in Mclath plage region 14179.
73
... ..
LI
~ ,,,,,, -~ - ,
interplanetary field line located at - 570 heliographic longitude for a solar
wind velocity of 400 km/sec. If the speed of the solar wind increases the interpla-
netary magnetic field line connection longitude is closer to the central reridian;if the speed decreases the connection longitude is closer to the west limb of the
sun. An illustration of this "most favorable propagation path" for arbitrarysolar wind speeds is shown in Figure 3.
Particle detectors on spacecraft at various locations in the interplanetarymediun usually record different particle intensities for tile same solar particleThe intensity is a function of the heliographic distance from the flare site and
the radial distance from the sun as the solar particles propagate through thesolar corona and out into the interplanetary medium along the magnetic field lines.A typical example is shown in Figures 4 and 5. Figure 4 gives the relative position
of three spacecraft during the solar particle event at 2147 UT on 30 April 1976.Hqelios I wns located at 1600 W, R = 0.65 AU, ilelios 2 was located at 156' W, R =
0.4 AU, and flIP 8 was located at 00 heliolongitude at 1 AU (i.e. orbiting the
earth. The solar flare that was associated with the solar proton event was locatedat 46 W. Thus tile longitudinal distance between these spacecraft and the flarelocation was 114 , 110 , and 46 for Helios I, Helios 2 and 11? 8 respectively
with the flare located to the east of the Helios spacecraft and to the west of the
IMP 7, 8
10 0.0. ,.-E 2
. ~ ~ I -I.4A _
0
LELIOSSU
C"I R - 0 4 AU
EElO
I- l- rr•4
j 14.O6SAU
R0 - 6 A
I MAY I?6 2 MAY IW
Figure 5. The 4 to 12 solar proton flux observed at three loca-tions in space; the earth (LMP 7 and 8) and at llo s I and Ile lt.s 2.
74
[IP S spacecraft. Figure 5 illustrates the tirk-intensity profiles of the solaIrproton Flux for this event as detected by sensors on these three spacecraft. Itis clear that a significant portion of the inner heliosphere was populated wit'isolar particles, with the detectors on IMP 8 responding to the larger flux aseould be expected from its more favorable location.
3. CONCEPT OF "SPACECRAFT EARTI"
Just -is a satellite can have detectors positioned at different locations onthe spacocraft, the earth has ground based detectors positioned at different loca-tions. Since the earth has a magnetic field, "spacecraft earth" can effectivlvlct as a satellite with its own built-in magnetic spectromaeter.
Solar particles travelling through the interplanetary mediun from the sun tothe earth are deflected by the geomagnetic field as soon as they enter the magneto-sphere. Their detection at the earth is primarily a function of particle energyand, for an anisotropic interplanetary particle distribution, a function ofdotector location. For detectors located above the atmosphere (i.e. balloonborne detectors or low altitude satellites) the only "shield" for these particlesis the geomagnetic field. For detectors located within the atmosphere, the verticalitnospheric mass of 1 1000 grams/cm
2 also acts as ar ibsorber. A ground-level
solar proton event implies that the solar particles nust have energies greater
Figure 6. Conceptual illustration of restricted asymptotic"ones of viewing sampling different sections of %pace.
75
than ~ 450 leV/nocleon, the minimum energy necessary to initiate a nuclear cascade
that can penetrate through the atmosphere and he recorded at sea level. Neutron
monitors located at sea level in high latitude locations (0 55 geomagnetic)
respond to particles greater than - 450 fteV; at lower latitudes the geomagnetic
cutoff is the principle factor controlling particle access. At geomagnetic latitudes
of - 45 , only particles with energies greater than - 3 GeV can be detected
by neutron monitors. In the equatorial region, the particles nust have an energy
of more than 1 L3 GeV to he detected. (For this paper we are considering only
particles incident in the vertical direction.)
The geomagnetic optics" or viewing direction of a neutron monitor is extremely
cotplicated. First introduced by Brunberg (1958) tl concept was refined into
asymptotic cones and directions by McCracken (1962) who advanced the technique by
the use of high speed digital computers. Figure 6 illustrates two types of asympto-
tic cones, one for a very high latitude (i.e. polar) station which has a relative ly"tight" cone viewing out of the ecliptic, and one for a moderately high latitude
station with geomagnetic latitude - 600 having a somewhat larger cone viewing in
the ecliptic plane. Stations at mid latitudes and equatorial locations have very
complex cones that can cover very large ranges in longitude. Figure 7 shows the
position of several asymptotic cones for high latitude stations. These cones have
been calculated using the trajectory-tracing technique (McCracken, 1962) with a
quiescent modeI of the geomagnetic field. Currents in the magnetosphere and the
magnetopheric tail were not included in the field model used to determine
these asymptotic cones.
Naturally, for a completely isotropic particle flux, knowledge of the geomagne-
tic optics, or viewing directions, of a station is not important, as the flux to
be detected is limited onlv bv the geomagnetic and/ or atmospheric cutoff. However,
most relativistic solar particle events are anisotropic, and although this makes
the analysis of these events more complicated, it does afford us the opportunity
to deduce some of the characteristics of particle propagation in the interplanetarymedium.
ISMPTOTIC LOMI4 DE
OLT
FIgure 7. Illustration of the asymptotic cones of acceptance
for selected high latitude neutron monitors as projected on
an extendcd world map.
76)
To analyze neutron monitor responses to anisotropic particle propagation, we
mast .ssuiTie a pitch angle distribution around the interplanetary magnetic fieldline. From our previous work (Smart et al., 1971; Smart et al., 1979) we assumethat the pitch angle distribution can be described by a Gaussian function with thewidth of the Gaussian curve controlling the degree of anisotropy. (See Figure 3.)
A detector "viewing" into the anisotropic particle flux as illustrated by cone Ain Figure 6 would record a higher increase than a detector viewing in a slightlydifforent direction, such as illustrated by cone B. Thus the concept that the"spacecraft Earth" can be used to deduce interplanetary propagation characterist icsfor relativistic solar particle events.
4. NEUTRON MONITOR RESPONSE
The response of a neutron monitor to an anisotropic flux of solar particles
can be calculated (see Shea et al., 1971) from the following equation:
I = Z S(Pi,X) J(Pi) F(Oi) A(P i ) APi (I)PC
where I is the relative intensity, P is the rigidity (Pc = cutoff rigidity),.S(P,X) Is the neutron monitor specific yield as a function of rigidity and
08 - - e-,1-
0
0 I -
o-
180 150 120 90 6C 0 30 60 90 120 1501 10
Figure t illustration of varios Gaussian pitch angledistributions around the interplanetary magnetic field line.
77
standardized depth (X) in the atmosphere (taken at sea level for this analysis),J(Pi) is the differential flux, in the interplanetary medium, of the relativisticsolar protons at each specific rigidity, PI, F(a i) is a function describingthe solar proton pitch angle distribution about the interplanetary magnetic fieldlines where a i is the pitch angle of the "'ith" particle, and A(P i) is aquantized function defining if a particle with rigidity P is allowed (A I) orforbidden (A
= 0) as determined by the trajectory calculations.
There is a unique set of values for the parameters defining the solar protondifferential rigidity spectrum, anisotropy, and apparent source direction that,when transmitted through the asymptotic cone of acceptance for each neutron monitorand through the neutron monitor specific yield function, will produce the observedincrease at any location on the earth. Using equation (I) the procedure isrepeated several times for different stations, pitch angle distributions, spectralslopes and direction of maximum particle flux in the interplanetary medium untilthe "best" agreement is found for the set of stations being considered. Usually,as the event proceeds in time, the anisotropy decreases and a new expression forthe Gaussian distribution must be used. In this way neutron monitor measurementsprovide an "instantaneous snapshot" of pitch angle distribution in the interplanet-
Thule
1h0l
Aooi ,Itler
o..oo A...
z
F-
2
> ,, roo T oy
Wa 6,,, Boyeore Adel~e4
W4 2 -2I 0'
MC Mul do Gossa
LI . . . . . . . lJ2100 2200 2300 UT 2MO 200 2300 UL
30 April 976 30 April 1976
Figure 9. The 30 April 1976 ground-level solar cosmicray event as observed by various neutron monitors.
78
ary mediUn for relativistic protons. It is emphasized that the anisotropies 4eare discussing in this paper are very large compared to those usually obs !rved byspacecraft.
5. THE RELATIVISTIC SOLAR PARTICLE EVENT OF 30 APRIL 1976 DURING STIP INTERVAL I
The ground level solar particle event of 30 April 1976 was associated with aflare at 2047 UT at S 08, W 46 (Solar Geophysical Data, 1976). The flare wasclassified as optical importance IB accompanied by major X-ray and radio emission.From an examination of the neutron monitor data from 10 locations, as shown inFigure 9, the earliest onset of relativistic protons appears to be between 2120and 2125 UT at Inuvik, Canada. This station also had the largest increase of 12percent. A comparison of the relative increases as recorded by the Inuvik andGoose Bay neutron monitors, is illustrated in Figure 10. Since both of these highlatitude neutron monitors are at "atmospheric cutoff" the difference in relativeintensity reflects the relativistic particle anisotropy present throughout thisu Veit.
Figure 7 shows that there is a 900 longitudinal separation between the asvmp-totic cone for Inuvik and the asymptotic cone for Goose Bay. Since Inuvik recordsthe earliest onset and the highest increase, we can conclude that the Inuvik asymp-totic cone of acceptance would be viewing into or close to the direction of maxinumparticle flux (i.e. along the interplanetary magnetic field line connecting theearth aith the sun). Satellite data place the interplanetary magnetic field lineat 21 N, 337* E in the GSE coordinate system (King, 1979). Translating to theearth's coordinate system shown in Figure 7, the interplanetary magnetic fieldwould intersect the earth at 200 N, 200* E longitude.
1N0 UVIK
-. 10
La- L, L, L~
fLa L. SAY
2100 2200 2300 UT
30 APRIL 1976
Figure 10. Relative increase observed by the Inuvik and GooseBay neutron monitors on 30 April [976. The difference in theobserved amplitude is an indication of the observed relativisticsolar proton flux anisotropy.
79
ASVNP 'OYIC LONG IMMi~
S •
T
Ii
Figure i1. Illustration of the pitch angles sampled by the var-ious narrow "asymptotic cone" neutron monitors. °The maximum rela-tivistic solar proton flux with pitch angle ( 303 is indicated bythe heavy shading. The light shading indicates the relativisticsolar proton fluxes with pitch angles between 30 and 60
Using this position as2the, direction of maximum flux and assuming a Gaussianfunction of the form et-(e /2.5)1 for the ptch angle distribution, w: findthat by using equation (1) we can obtain calculated neutron monitor increases ingood agreessent with the observational data. Figure II is similar to Figure 7except that the shading indicates the geocentric directions n Interplanetaryspace that contain the most intense particle fluxes. The heavy shading representsthe geocentric radial direction containing particle fluxes with pitch angles 30
°
oroess; thle lighter shading represents particle fluxes with pitch angles between30 and 60}
in comparing this figuire with the increases in Figure 9, we nute that thelargest increases should he recorded by the louvik and Tixie Bay neutron monitors,as indeed is the case. (The Kerguelen Island asymptotic cone is not illustratedin Figure II as it covers a wide range of longitudes. This cone, illustrated inShea and Smart (1975), "sweeps" across this area of maximutm intensity consistentwith the increase shown in Figure 9.) In contrast, the Alert, Thule, Goose Bayand Mc'lurdo asymptotic cones are far removed from the direction of maximum particleflux - in keeping with the much smaller increases recorded at these stations.
Considering stations looking in the forward direction such as Inuvik, andcomparing them with stations looking in the reverse direction such as Goose Bay,ae can see that the "forward to reverse" particle flux ratio is at least six toone, Determining a Gaussian distribution that when coupled with the asym~ptoticviewing directions and neutron monitor spectrum/yield function produce calculatedincreases close to the observation data results in a function of el-(B /2.5)].Since these asymptotic cones actually sample a range of pitch angles we will assumethat the forward flux contains pitch angles between 00 and 600 and that the reverseflux contains pitch angles between 120 , and 180 . Evaluating our Gaussian formbetween these limits we find that the average "forward to reverse" flux ratio isabout a factor of 10. Figure 12 illustrates the probable pitch angle distributionof the relativistic particle flux in interplanetary space and also illustratesthe region of the forward flux and backward flux.
80
..... .~~~ i $I.. .OWA.........
30 APRIL 1976
2145 UT
I--
IJ
W 0.5 i7~'7',/-I777/'i
/ '/ ,FORWA R D/ ILUX/ '
30 60 90 120 150 180
ANGLE FROM INTERPLANETARY MAGNETIC FIELD
Figure 12. Derived relativistic solar proton pitch angledistribution during the maximum of the 30 April 1976 ground-
level solar cosmic ray event.
6. TIlE RELATIVISTIC SOLAR PARTICLE EVENT OF 22 NOVEMBER 1977 DURING STIPINTERVAL IV
The relativistic solar partaicle event of 22 November 1977 was associated witha flare of importance 2B at 0945 UT. This flare was located at 240 N, 400 W (SolarGeophysical Data, 1978), again near the favorable propagation path between the sunand the earth. Since neither solar wind nor interplanetary magnetic field dataare available from U.S. spacecraft, we assume a direction of maximum particle fluxalong an interplanetary magnetic field line intersecting the earth at a longitudeof 3490 (i.e. 9 AM local time) for the onset of the event and at a longitude of330' for the time near maximum Intensity.
Figure 13 illustrates the increase in cosmic ray intensity as recorded bysevva neutron monitors (Komori, 1977). This was also an anisotropic particleevent although the degree of anisotropy was much less than for the event on 30April 1976. By placing the direction of maximum particle flux (longitude 3300 E)on the nlap in Figure 7, it is evident that the Goose Bay and South Pole neutronaonitors should have detected the largest increase with the Inuvik monitor, whoseasymptotic cone is - 1200 away from the longitude of the maximum particle flux,detecting a much smaller increase. ije must inject a note of caution here as thereis a large altitude gradient in the atmospheric secondary neutron flux and theSouth Pole monitor is located at an altitude of about 3000 meters. This altitudecorrection has not been applied to the South Pole data shown In Figure 13.
For this event we find hat a particle pitch angle distribution described bya Gaussian of the form el-(O /8.0)) results in the best fit to the experimentaldata. Again taking the range of pitch ang' ms between 0' and 600 for the forwardflux and between 1200 and 180* for the reverse flux and evaluating the Gaussianexpression between these limits we find that the average "forward to reverse" flux
81
I GLE on 1040 UT Ntov. 22. 1977
am0 - ALERT 12000 -
19 1
7000 11000ow- 1040 UT
90W- 220-
8000 - DEEP RIVER 20-
26 %
7000 1200 - 1S0-
9000- 1100.. PMCPIR002S
GOOSE BAY pUR8000 - 311 1000 -
j 7SE"SAY1042 VT
1040 UT 1700 -
7 8 0 - S O U T H P O E I WS29 11Is -
7400 - IJI 30
7000 -
20 FLARE PTARTED 094e5 LIT ENDED 1105 UT NOV. 22R23 W40
Figure 13. The 22 November 1977 ground-level solar cosmicray event as observed by seven neutron monitors (afterKomori, 1979).
ratio is about a factor of 2. An illustration of the Gaussi~n function for thisevent and the forward and backward particle flux is shown in Figure 14.
7. CONCLUSION
We have summarized the neutron nonitor data for the relativistic solar parti-clo events that occurred on 30 April 1976 and 22 November 1977. Both of the solarflares that produced these relativistic particles occurred in a location on the sun(46" W4 and 40" W4) favorable for maximum intensity to be recorded at the earth.
82
22 NOVEMBER 19771100 UT
1.0
z
I-
l- z FACKWARD,'/
W
0.0 11 , , I I I. . . . .r i .
30 60 90 120 150 180
ANGLE FROM INTERPLANETARY MAGNETIC FIELD
Figure 14. Derived relativistic solar proton pitch angle
distribution during the maximum of the 22 November 1977ground-level solar cosmic ray event.
From this study of these two relativistic solar particle events using theearth as a detector in space (i.e. satellite Earth) we can determine the pitchangle distribution in space throughout these events. Although there was an aniso-tropic particle distribution for both of these events, the degree of anisotropy wasconsiderably different, the relativistic particle flux "forward to backward" ratiobeing 10 to 1 for the event of 30 April 1976 and 2 to I for the event of 22 Novem-ber 1977.
References
Brunberg, Ernst-Ake: 1958, "The optics of cosmic ray telescopes", Arkiv for Fysik,14, 195.
King, Joseph, H.: 1979, Interplanetary Medium Data Book Supplement 1, 1975-1978,Report No. NSSDC/WDC-A-R&S 79-08, National Aeronautics and Space Administra-tion, Goddard Space Flight Center, Greenbelt, Md.
Komor, 11.: 1979, "An analysis of relativistic cosmic ray increase on 22nd November1977", 16th International Cosmic Ray Conference, Conference Papers, 4, 19.
HcCracken, K.G. : 1962,"The cosmic ray flare effect. 1. Some new methods of anal-
ysis", .. Geophys. Res., 67, 423.
Shea, M.A., and Smart, D.F.; 1975, "Asymptotic Directions and Vertical CutoffRigidities for Selected Cosmic-Ray Stations as Calculated Using theinternational Geomagnetic Reference Field Model Appropriate for Epoch 1975.0",Report No. AFCRL-TR-75-0247, Air Force Cambridge Research Laboratories,
Bedford, Mass.
83
References.vakart , 1). F., Shi, M.A., liunblIt!, J. K. , an d TanskaRenR, P.1. 197 1,
"A determination of the spectra, spat il anisotropy, and prop-aation characteristics of the relativistic solar cosmic-
ray flux onl November 18, 1968", 12th International Conferenceoii Cosmic Rays, Hobart, Conference Papers, 2, 483.
SnrI. F., Shea, MI.A. , Humhle, J.E. , and Tanskanen, P.J.: 1979,"A\ model of the 7 May 1978 solar cosmic ray event", 16thInternational Cosmic Ray Conference, Conference Papers, 5,2 33.
Solar Geophysical Data: 1976, U.S Department of Comme~rce, Boulder,Colorado, No. 336.
Solar Ceophysic-al Data: 1978, U.S. Department of Comnierce, Boulder,Colorado, No. 405.
Discuss ion
McLean: Could the special characteristic of flares which determines whether theyproduce GLE, be proximity to thke magnetic field line which connects to theEarth? The two examples which YOU gave uis fit this idea.
Shea : The shape of the 'time-intensity" profile of relativistic solar particleevents as detected at the Earth by neutron monitors appears to he a function ofthe heliocentric distance between the associated flare and the most favourablepropagation path, defined as the "idealized" magnetic field line connecting theEarth with the Sun. However, flares associated with ground-level events haveoccurred at many solar longitudes. Three ground-level events, those on20 November 1960, 30 March 1969 and 28 January 1967, are associated with flareson the invisible part of the solar disk (i.e. 90'%)
Dryer: Is it possible to usO the expanding shock front (temnporally and spatial lv)as a particle mTOdUlat iOn eechlanisr. in the cast- of th, diifcren t Hel ios 1 and 2temporal profiles whenl they were both ,100 west of the flare sito'?
Shea : III Our Opinlion itWould be' \erv dif fi colt in view of the time associations.In interplanetary space thle radially propagating shotcks often act as "boundaries"between different part i cle popuilat ions; however, in this case yVeIl would requiretile equivalent oI a More ton wave to propagate very rapidly (i.e. orders ofmagnituode fast ,r tha1n the' observed veoiCfty) to 1100 from the flare site and thensuddenly act as ;I "quasi-stable" harrier between thke two satel lites with nofurther l ateral p ropagat ion. l1c sc ient ists with detec tors on the Hl~dios I and 2space p~robes havc rt*ported that- there applears to be a taogeut 131 discont inuityaSSOC iated with a to-ro tat ing st ream between thle two ipatvcratit and we feel thisi5 a Much MoreO plausible explanation.
84
8. Thermal Cyclotron RadiationFrom Solar Active Regions
V. V. Zheleznyakov and E. Ya. ZlotnikInstitute of Applied Physics, Academy of Sciences,
Gorky, U.S.S.R.
Abstract
Spectra with fine structure resulting from the thermal cyclotron radioemission from solar active regions are discussed. Source conditions (distributionof magnetic field and kinetic temperature with height) are suggested which wouldcause the spectrum to take the form of a set of cyclotron lines and high-frequencycut-offs. Each distribution results in a characteristic spectrum and polarization.This permits one to deduce the conditions in the source through the properties ofthe observed microwave solar radio emission. To obtain reliable data on the finestructure and judge conditions in the sources it is necessary to study microwavesolar radio emission with swept-frequency or multi-channel receivers and highlydirectional antennae.
85
1. INTRODUCTION
All the phenomena in microwave solar radio emission occur against theS-component - slowly varying radiation of the local sources above sunspots andflocculi in the active regions of the orona and chromosphere. The S-component ischaracterized by a smooth frequency spectrum with the maxima in the wavelengthrange 5-10 cm, unpolarized at longer wavelengths and polarized in the extraordinarysense at shorter wavelengths. These peculiarities were explained in terms ofthermal cyclotron and free-free radiation mechanisms by Zheleznyakov (1962, 1963)and Kakinuma and Swarup (1962). Further progress in the theory of the S-componentwas made by analysing various models of active regions, calculating the expectedcharacteristics of cyclotron (and free-free) radio emission in these models, andcomparing these calculated characteristics with the observed ones (see Livshitset al., 1966; Ziotnik, 19b8; Lantos, 1968). However, Zheleznyakov and Zlotnik(1979, 1980a,c) paid attention to the fact that thermal cyclotron radio emission
in neutral current sheets of the corona causes a set of discrete cyclotron lines.If the kinetic temperature of a plasma in the sheets is sufficiently great(T - 107_108 K) these lines may be observed against the S-component mostly as radia-tion at the second and the third harmonics of the electron gyrofrequency WB
= eB/mc.
Fine structure such as cyclotron lines should be expected in the pre-flare phasesof developing active regions when the formation of hot neutral current sheets ism.-st probable. This possibility of narrow-band features in the spectrum of micro-wave radio emission and of cut-offs and narrow-band features in the spectrum ofradio emission of solar local sources (Kaverin et al., 1976) stimulated the presentauthors to consider the fine structure more thoroughly. The results of theanalysis are given below.
2. RFNiE% OF BASIC TiEORY OF S-COMPONENT
Before discussing various forms of the fine structure in the microwave solarradio spectrum, let us recall the peculiarities of the cyclotron mechanism in theinhomogeneous magnetic fields above sunspots and the conditions of the typicalS-component (see in more detail Section 13 of Zheleznyakov (1977)).
The non-relati'istic electrons in a magnetic field are known to radiate andabsorb the electromagnetic waves at the frequencies equal to or a multiple of thegyrofrequency a - sR where s = 1, 2, 3 .... Therefore cyclotron radiation andwave absorption at a given frequency in the inhomogeneous magnetic field occur inthose gyroresonance layers where the gyrofrequency is
W B :z "I/s -(1)
Under typical conditions of the solar corona and chromosphere the optical
thickness of these layers T is different for the ordinary and the extraordinarywaves. For the ordinary waves the optically thick layers 0 ' 1) correspond tothe harmonics s < 2; for the extraordinary waves such layers correspond to thcvalues s < 3. In both cases the gyroresonance layers are optically thin (1 1 1)for the higher harmonics.
86
A theory of the S-component is based on a model of the active region above
a sunspot represented in Figure 1(a). According to Eq. (1), the gyrofrequency WB(in a magnetic field B) decreases monotonically with height h over the photosphere.
TB
C
h)hol
Figure 1. The model of a standard source ofthe S-component: a) distribution of themagnetic field B and the kinetic temperature T
over the height h (h 0 is the lower edge of thecorona, B0 is the magnetic field at this level);b) the frequency spectrum of radiation; solid
and dashed lines here and elsewhere correspondto extraordinary and ordinary waves.
The kinetic temperature of a plasma increases from the chromospheric values(Tch - 104 to 3 x 1
04 K) to those typical of the active corona (Tc - 3 x 106 Q.
In such magnetic fields the gyroresonance layers (Eq. 1) corresponding to thehigher harmonics are situated at greater heights in toe solar atmosphere (closerto the observer). This fact, along with the above value3 or optical thickness for
the ordinary and extraordinary waves of gyroresonance layers, leads us to concludethat the second gyroresonance layer radiates the ordinary, and the third gyro-resonance layer radiates the extraordinary waves. The brightness temperature ofthe ordinary component of radiation is
Tb T(h2) , (2)
87
whlere h:_ is the height of tile second gyroresonance layer with the gyrofrequency/;thle brightness tenmperature of the extraordinarv mode is
T--T(h 3 ) , (3)
where h-. is thle height of thle third gyroresonance layer -' 3 (see Eq. 1).Radiaotion from thle deeper ladyers is absorbed at the heights If: and h1 for theoltd marv and ext raordinary modes respect iveI rv he con trihbution of the opticallythin hi gher harmionics is insignificant. Ihe above argument enables us to derivethe spe'tr tom of thle S-component, Thi..) , for both raves. Such a derived spectrumis given in Figure 1(b). From this figure it is clear that for the extraordinaryraIdi it ionl 1, =. a, t thle frequencies . <
3,,,B,, corresponding to thle position of thle
third ;vroresoilance level Inl tife Corona - at~'the heights It IbI At higher fre-11uetll"ch's tlt, level S 3 sinks to the chromosphere. In -. e transient region Tbdiops, tending to the valuLe TLh. The same is observed tot the ordinary mode hottill dro'p Inl h begins earlier -at the frequency corresponding to the mnoment oftilt' (ItrIn"it ionI Of thel level '. 2 to tih jegdts h) -: J) From Figure 1(b) itfollows that tile thermal cyclotron radiation in a tandard model of thle S-componentIS uopolirized at thle low t.- , 2 I.j and the higher I.. 3-,B) frequencies and israther strongly polarized in the ranige of the intermediate trequencies2.I -Bo f , B or wnich one or bothl of the gyroresonance layers S = 2,3 ark,Sit lte
1L in tilt? ttzosi tion region hetween thet cittomosplierc. and tile corona where
the, temperiture gradient is strong.
ishlold be note-d that in the, present report all thle frequency spectra arelrlpr~~=illIt, Itttl I) rathler than ill specific initensity 11,). 1,o tranls-
trm 1;( into dlpetidinu- ,I 1 onf one shlould take into accoulnt thlat 1 BA's 11 T I t t )10 SPC t Itonf I .( . ) has a tiaXttill at the frequency Ma \ i 1 andit bespec I tr;i'''c' r
0e onar- Litl J~ 4 it io o tIf tie S- componenlt Ita s .I1a lxm 1lm a t t ?I
t 1' 1' .Il '
III , ,i. ,I tie treqlitto'%sp5ectrum'. wd pIariz~it ion of the S-I ~ttl't I~ ~ . ,fi i r7-1i,, bc\ OhL rVA te jlt1 )I tI. lrllItOW'.Ve a d " 1 2, e i ss i ) 1
1!11 11 I. ls 0IL1 t12IV V0,ltS 1.1freicll et al. 19-01. It is cleatIOl'l'et 'illtt :i ,ollsid, red doe's not exhiaus t .11 1 possibl 11on)tdi t i 1nIS in I) L
'I"t I ye eI '. L! I ')n -1,1L'(1 111 .V t Ite s p e t ra If r ad io, 7 e i ss ion mav' ;e tar More
:I. X N1OPELATI. i .M %I %\I' I N1 (F , %XGM TIC FIELD) TREN(;T1 IN TIlE CORON k
let Is5 IIs Ill> t aSple, t IU n ,t the-rinalI vc lotro tr alli at ion in a fiodel where~ttilt' vaIlueof 't ie maIgnet ic f ieldk reacihes its maximutm at a cettainl height hma fortilt' previous list tibllt iotn of temporatulre (see Figure 2(li ). For example, sucha1 list rib It ion ofl thll magnet tIc f ield Can be' rte I i .7t( in a gro~ip of sunspOts repre-se nted Ini Figutre !i.i ) -'vt Il isregarlling thte elect ric culrrents in the coronial andtilt- llromos-pierI c plaisma ( Swee(t, 195,8). (Here 3lIttg tile axis It passing thtroulgh tite11,111 rt I -r 0 I li, ,t tilt. magne-tic field tlle masimllm h. iso tetainlv present. if
"fill It whet.' If is the( lo)wer vdlgc of the' lowi, i brti the model considered isIll r-dl t,' lc'l t lt, sI indairJ mode I. 11,11 i f I it. then tilt- radialtion
A I* is im 711 id h.i i F % j gotrt .1 i , t il' iI t i gh t Ti 5
ss tll i a , t i'I dt ro I, I oWa tds tile
li tjwr t 1 -1. ton i 's 4t thI I p intts -B . ,- it Ti. i r ittit'I wa',' -s ) aoB
--. rw tli~ iv ~i,, ) jj a on t I jis ) h, raw cti, I tI s t. 18
AD-A30 687 PROCEED ROS OF THE TP (STUDY OF TRAVELLNG 213IERPLANETARY PHENOMENA SY.U AIAR FORCE GEOPHYSICSLAB HANSCOM AFB MA M A SHEA ET AL 27 DEC 82
UNCLASSIFED ADGL-R-82-0390 FG32 NEhEmhEImimhImEIIEIIIIIIIEIIEIIIIIIhEEEII IIIIIIIIIIIIlI IIIIIIIIIIIIlIIIIlfllfflfflffllfllfIIIIIIIII~lEEI
_____ 2.28 11112.21L. IW mm1 .
MICROCOPY RESOLUTION TEST CHARTNAIOA BUREAU Of STANDARDS -963 -0
T B a)
ho h ,mcx
r
TC
2 max 3 mx max
C)Figure 2. The model of a source with the maximumof the magnetic field over height: a) magnetic linesof force above a group of sunspots (hmax is the
height at which the magnetic field has its maximumvalue); b) distribution of the magnetic field andkinetic temperature over the height; c) the fre-quency spectrum of radiation.
values exceeding Bmax . Note that such values may exist at small heights in thechromosphere below the neutral point. In this case however the kinet.c temperatureis low; thus the contribution of these regions to radiation is neglecteu. in thefrequency range
2(jB < w < 3uB the radio emission is strongly polarized (withSmax
an excess of the exTaordinary wave). The degree of polarization approximates100%. Moreover, it is important to note that at the frequencies 3,,1B x
4',BxIl ma ° "
the cyclotron lines emerge because of the significant increase of opticalthickness of gyroresonance layers caused by a small change in the magnetic field
89
at the heights h z hmax. For the optically thin gyroresonance layers the increaseof optical thickness leads to a corresponding increase in the brightness tempera-ture at the frequencies w - sw . The radiation in the line w c
3Bm consists
only of the ordinary waves. The lines at the higher harmonics are partly polar-ized with an excess of the extraordinary waves.
The distinctive feature of the model in Figure 2(b) with the magnetic fieldmaximum at a certain height in the corona is the fact that the fine structure inthe spectrum (a set of cyclotron lines and high-frequency cut-offs) appear underthe condition where sharp gradients of the magnetic field and electron density areabsent in the corona and chromosphere.
The calculated profiles of the cyclotron lines and high-frequency cut-offs arerepresented in Figure 3 (for the parabolic approximation of the dependence B(h)near the maximum of the magnetic field).
)
-4-2 02Figure 3. The calculated dependence of Tb in thecyclotron line a) and in the high-frequency cut-off
b) on the parameter z = (w-swB ax)// -m T cos a inthe model with the maximum of tMe magnetic field(Figure 2(a)).
A, the harmonics s for which the optical thickness is small at the frequencyw= B a we obtain a line the form of which is represented by the curve a) inFigure 'a? The halfwidth of the line Aw is determined by the Doppler effect
Awcos 0E (4)
90
where 3T = VT/C, VT is the thermal electron velocity, a is the angle between the
magnetic lines of force and the line of sight.
In the case where the optical thickness at these frequencies is well aboveunity, the frequency spectrum takes the form of a high-frequency cut-off given bythe curve b) in Figure 3. The characteristic width of the cut-off Aw is aspreviously determined by relation (4).
4. RADIATION FROM A THIN MAGNETIC TUBE, FILLED WITH HOT ELECTRONS
Judging from the X-ray emission data (Cheng, 1977) it is possible for the
force tube of a bipolar magnetic field above a sunspot to be filled by hot electrons(Figure 4(a)). These electrons may, in principle, fill the whole of the tube or be
present as an admixture to the electron component at normal coronal temperature.
As an example we treat the case where a relatively thin force tube is filled
by hot electrons (with T )0 T). Let us also assume that their pressure is smallcompared with the magnetic pressure and that therefore the pressure of hot electronsdoes not markedly change the value of a magnetic field in the tube. The assumeddistribution of T and B along the axis h is shown in Figure 4(b).
It is easy to see that in the framework of this model the frequency spectrumof cyclotron radiation escaping the corona along the axis h will represent a super-position of the usual spectrum of the S-component (Figure l(b)) and cyclotronlines at frequencies which are multiples of ()Bt (wBt is the gyrofrequency ina force tube at the height ht). This spectrum is given in Figure 5(b). Analysis
shows that the force tube can radiate only the cyclotron lines w z 2wBt, 3WBt .....
The second line is totally polarized in the sense of the ordinary wave. The linesbeginning with the third are partly polarized, with an excess of the extraordinarywa ve.
5. RADIATION FROM A HOT NEUTRAL SHEET
It has been mentioned at the beginning of the report that the fine structureof a spectrum may be due to the cyclotron radiation of hot neutral current sheets(Zheleznyakov and Zlotnik, 1979, 1980c). Probably these appear in the lower
corona or in the upper chromosphere during the pre-flare phase of solar activity.Thus detection of the spectral fine structure associated with neutral currentsheets may prove to be an important method of predicting and recognizing suchsheets.
The distribution of the magnetic field and temperature in the corona alongthe line of sight passing through a neutral current sheet is given in Figure 5(a).It is assumed here that the plasma has an enhanced temperature only within thelimits of the neutral current sheet, including the edge regions with a magneticfield B < B0 .
The frequency spectrum of thermal cyclotron radiation from a neutral current
sheet approaching the point ho is seen in Figure 5(b). It represents a set ofrather broad lines at the harmonics of the gyrofrequency w = su%0. The calculation
91
hh
ht h
C )
Figure 4. The model of a source with hot electrons fillingthe magnetic field force tube: a) magnetic lines of forcewith hot electrons (close hatching) and those crossed bythe antenna pattern (wide hatching); b) distribution ofthe magnetic field and temperature over the height (Bt Isthe magnetic fie d at the level ht of the force tube);c) the frequen.', spectrum of radiation.
92
T B
T ----- T-
-(Iii' I'Wo 2wco 3wj OwO
W B 2w 3ca 4w 0 fW
C)
Figure 5. The model of a source with a neutralcurrent sheet and a monotonically vanishing magneticfield: a) distribution of the magnetic field andkinetic temperature over the height (ho is theconventional edge of the neutral current sheet);b) the frequency spectrum of the proper radiation ofthe neutral current sheet; C) the frequencyspectrum with allowance for absorption and radiationin the corona.
93
of such a spectrum has been discussed in ample detail by Zheleznyakov and Zlotnik(1980a) so it is omitted here. Note only that the appearance of discrete cyclotronlines is due to the generation of radiation in the region of a quasi-uniform
magnetic field B0 at the edge of the neutral sheet. The source is restricted frombelow by the level at which the plasma frequency w is equal to the harmonicfrequency. It should be readily apparent that the given level in the non-relativistic plasma must be situated far from the middle (i.e. at tile edge) of thesheet. In fact, from the condition of static equilibrium of a plasma in the
magnetic field of a neutral sheet B2/8c = 2N0 K T it follows that oP = wB /2wT(No and 10p1 are the density and the plasma frequency in the middle
0of th8 sheet).
Since tile parameter 6T < 1, the last equality implies that -)P >> UwB . At low
harmonics of the gyrofrequency wi = sw B this inequality ensures the validity of thecondition wo > w. Hence it is evident that the plasma level P = W for the
emitted trequencies is located far from the middle of the sheet in the region wh re
the magnetic field is close to B0 .
When escaping a neutral current sheet the cyclotron radiation does not containthe extraordinary mode at the frequency '-)B0 , since n
2 -- 0 in the sheet. The
lines w = 2w 8 0 , B " are unpolarized or partly polarized with an excess ofthe extraordinary wave depending on whether the source is optically thick oroptically thin. The spectrum and the polarization of radiation escaping thecorona certainly differ from those in Figure 5(b) owing to absorption of the
neutral current radiation in the upper regions of the corona and of the propercyclotron radiation of these regions. For example, the line Uo = coB1 in the
ordinary radiation and the extraordinary mode in the line uo =2,t0 will vanishowing to absorption in the corona (the first will vanish at the gyrorcsonancelayers , = 2(,t, and the second at the level wi = 3LB). In general the observedspectrum takes the form presented schematically in Figure 5(c). The high-frequency cut-offs at the frequencies .,) = 2wB0 and " = 3
, 0 . are associated withthe limiting from below of the source of the S-component in the corona: tileradiation from the deeper regions of the corona below the neutral current sheetdoes not penetrate through the sheet at the frequencies wo U p (on the emissionpropagating through neutral current sheets in the solar corona see Zheleznyakovand Zlotnik (1980b)).
Let us stress once again that thermal cyclotron radiation is produced in tileneutral current sheet by only the external regions of the sheet. The behaviour of
the magnetic field toward the deeper layers (drop to zero and reversal of sign orreducing to a certain limiting value followed again by increasing) does not affectthe formation of the frequency spectrum as a set of resolved cyclotron lines.
Therefore the radiation with the characteristics shown in Figure 5(c) may, inprinciple, be generated by hot electrons trapped by the local minimum of themagnetic field in a force tube. This minimum may occur in a force tube if the
energy density of hot electrons is comparable with that of the magnetic field.
6. COMPARISON OF MOI)ELS
All our models of magnetic field and electron density distribution with
height give results differing from the standard model of the S-component. They
result in frequency spectra with fine structure in the form of cyclotron lines and
high-frequency cut-offs. Every model however has its own characteristic spectrum
and polarization of thermal cyclotron radiation. This circumstance permits us todeduce the distributions N(h) and T(h) from the observed spectrum and the polariza-tion of microwave radiation of solar active regions. From the above it is clear
that for the detection of cyclotron lines and the study of fine structure,
94
simultaneous spectral and polarization measurements in the wide frequency range(covering at least a factor of three in frequency) and with the adequate resolutionAw/w ' 0.05 are necessary. The spectrographs should be combined with the antennasystems with as much directivity as possible: beaswidth must be less than a sun-spot or the distance between sunspots in a group. Antennae similar to the RATAN-600 or the 100-m Bonn radio telescope are preferable, otherwise wien receiving theradiation from the whole active region the fine structure is smoothed and one canhardly detect with confidence the cyclotron lines and cut-offs; their identifica-tion and interpretation then become doubtful. An example of this smoothing isgiven by the radiation from a magnetic force tube (see Figure 4). If the trapped
hot electrons are distributed along the tube over the region where the magneticfield varies greatly (AB/B _ 1), then in the total radiation the cyclotron lineswill diffuse and the spectrum will be smoothed. Conversely, if the receiving Isperformed by the antenna with narrow knife or pencil directivity pattern (thecross-section of this pattern is represented in Figure 4(a)), then the frequency
spectrum of the observed cyclotron radiation will contain discrete cyclotron lines.
A detailed discussion of the problems tackled in the present report is
offered in papers by Zheleznyakov and Zlotnik (1979, 1980a,c).
7. CONCLUSION
For many years solar microwave flux and polarization have been measured ata few discrete well-spaced frequencies. The present study suggests that itobservations could be made with wide-band spectrographs and spectropolarimetersusing highly directional antennae, then valuable information on cvclotron linesand spectral cut-offs would be obtained. These in turn would give valuableinformation on the magnetic fields and temperatures in solar active regions.
References
Cheng, C.-C.: 1977, "Evolution of the high-temperature plasma in the 15 June 1971flare", Solar Phys., 55, 411.
(;el'freich, G.B., Akchmedov, Sh.B., Borovik, V.N., Gol'nev, B.Ya., Korzhavin, A.N.,
Nagnibeda, B.C., and Peterova, N.G.: 1970, "A study of local sources of theslowly-varying component of solar radio emission in the cm-wavelengths range",Izv. Glavonoya Astron. Obs., No. 185, 167.
Kakinuma, T., and Swarup, G.: 1962, "A model for the sources of the slowly varying
component of microwave solar radiation", Astrophys. J., 136, 975.
Kaverin, N.S., Kobrin, M.M., Korshunov, A.l., Tikhomirov, V.A., and Tikhomirov,Yu.V.: 1976, "Fine structure in the spectrum of the S component of solar radioemission at 4.5-7 GHz", Pls'ma Astron. Zh., 2, 577; Engl. transl., Soy. Astron.Lett., 2, 229 (1976).
Lantos, P.: 1968, "A model for thermal gyromagnetic radio emission from solar
active regions", Ann. Astrophys., 31, 105.
95
Referencos
Livshits, N.A., Obridko, B.H., and Pikel'ner, S.B.: 1966, "Radio emission andatmospheric structure above sunspots", Astron. Zh., 43, 1135; Engl. transl.,Soy. Astron. A3, 10, 909 (1967).
Sweet, P.A.: 1958, "The neutral point theory of solar flares", in Electromagnetic
Phenomena in Cosmical Physics (ed. B. Lehnert) (Proc. lAU Symp. No. 6), p. 123(Cambridge Univ. Press).
Zheleznyakov, V.V.: 1962, "The origin of the slowly varying component of solarradio emission", Astron. Zh., 39, 5; Engl. transl., Soy. Astron. AJ, 6, 3(1962).
Zheleznyakov, V.V.: 1963, "The frequency spectrum of the slowly varying component
of solar radio emission", Astron. Zh., 40, 829; Engl. transl., Soy. Astron. AJ,7, 630 (1964).
Zheleznyakov, V.V.: 1977, Electromagnetic Waves in Cosmic Plasma, Nauka, Moscow.
Zheleznyakov, V.V., and Zlotnik, E.Ya.: 1979, "Diagnostics of neutral currentlayers under space conditions", Usp. Fiz. Nauk., 128, 725; Engl. transl.,Soy. Phys. Usp., 22(8), 665 (1979).
Zheleznyakov, V.V., and Zlotnik, E.Ya.: 1980a, "Thermal cyclotron radio emissionof neutral current sheets in the solar corona", Solar Phys., 68, 317.
Zheleznyakov, V.V., and Zlotnik, E.Ya.: 1980b, "Fine structure of microwave solarradio emission from solar activity centers", Astron. Zh., 57, 778; Engl. transl.,Soy. Astron., 24(4), 448.
Zheleznyakov, V.V., and Zlotnik, E.Ya.: 1980c, "Influence of neutral current sheets
in the cosmic plasma on the frequency spectrum of transmitted radio emission",Astron. Zh., 57, 1038; Engl. transI., Soy. Astron., 24(5), 595.
Zlotnik, E.Ya.: 1968, "Theory of the slowly changing component of solar radioemission. I, II", Astron. Zh., 45, 310, 585; Engl. transl., Soy. Astron. AJ,
12, 245, 464 (1968).
Discussion
Alurkar: What's the width of the line emission you are expecting?
Zlotnik: The width of the line is determined by the Doppler effect and so must be
of the order of AT - the ratio of the electron thermal velocity to the light
velocity. If the kinetic temperature is of the order of 108 K, then Aw/. - 0.1.
The line may be more narrow for lower temperatures.
Schwenn: Is there any relation between your results and Dr. Alurkar's measurements
shown this morning?Zlotnik: I'm not sure, because my results are for centimeter-decimeter wavelengths
rather than decametric wavelengths which Dr. Alurkar is concerned with.
96
W0
0N
9. Observations of the Quiet Sunat Decametre and Metre
Wavelengths
A. A. Abranin and L. L. BazeijanInstitute of Radiophysics and Electronics,
Academy of Science, Ukrainian S.S.R., Kharkhov, U.S.S.R.
R. A. DuncanDivision of Radiophysics, CSIRO,
P.O. Box 76, Epping, N.S.W. 2121, Australia
Abstract
During the STIP interval, April 1976, the very low level of sporadic radioemission of the Sun at decametre and metre wavelengths presented us with a goodopportunity to observe the radio emission of the quiet Sun.
1. DECAMETRIC OBSERVATIONS
Decametric investigations of the radio emission of the quiet Sun are diffi-cult because of the weakness of the radiation and the high brightness temperatureof the cosmic background. These difficulties are increased by the ionosphere andby strong interference from radio stations. The ionosphere produces undesirableattenuation and fluctuations of signal because of ionospheric inhomogeneities. Asa result even the flux density of the quiet Sun for frequencies <30 MHz is hardlyknown (Cronyn and Erickson, 1965; Bazelyan et al., 1970; Aubier et al., 1971).Of the structure of the <30 MHz quiet Sun nothing is known.
97
In April 1976 this kind of measurement was attempted at 25 MHz at Kharkov.The period was one of very low radio emission activity. As well as two-dimensionalmaps of the Sun obtained with a large radio telescope, observations of the partialsolar eclipse of April 29 were made. Because of interference however only thefirst part of the eclipse was successfully observed. We present here a shortdescription of the equipment and some of the observational results.
The observations were made using the radio telescope UTR-2, which has fivesimultaneous beams separated in declination by 24' arc (Braude et al., 1971). Eachbeam has a half-power width of 28' arc in right ascension and 35' arc in declina-tion. Phase-modulated radiometers were used as receivers. The bandwidth of thei.f. amplifiers was 5 kHz, and the time constant was 15 s. To reduce the effect ofinterference two independent receiving channels were used, separated in frequencyby 50-100 kHz. Observations of the Sun were made daily at ±2 b from local noon.They consisted of several series of records, each series consisting of three scans,although, in principle, one scan is enough. The declination of the central beamfor the three scans was
60 -0'.2, 60 and 6( +0.2. The duration of each scan was
10 min.
The first step in processing was scaling the scans using the calibrationsignal of the radiometers and the parameters of the radio telescope. Then therecords of the two independent channels were averaged and smoothed over an intervalof 2 min. This interval corresponds approximately to the half-power width of thebeam. From these smoothed scans radio maps of the Sun were constructed.
The resultant maps of the brightness distribution of the quiet Sun forparticular times are given in Figures la, lb. One is for one hour before localnoon and the other for one hour after local noon ( 8 h 2 5 m and ]0h)5m II'). Tleoptical disk of the Sun and its active regions (Solar Geophyiscal Data, 1976a), andthe corona at 5303 X (Braude et al., 1971) are also shown. The coordinates aredeclination in degrees and "right ascension" in minutes of time measured from thecentre of the Sun.
During the period of observation and the preceding days, solar activity waslimited to the Mclath regions 14179 and 14185. It was mainly weak chromosphericflares and subflares (Solar Geophysical Data, 1976a). Two centres of activity were
observed in the radio wavelength range 8.5 mm to 2 cm also (Solar Geophysical Data,1976a), but at decametre wavelengths these two regions are indistinguishablebecause of either low intensity or small separation.
Figure I shows a very large shift of the radio image of tile Sun from itsoptical position. As the shift of the radio image is in the same direction fortwo practically symmetrical observations either side of noon, interpretation interms of ionospheric refraction is unlikely. Furthermore, on the preceding day,1976 April 29 (Figure 2a) the maximum of radio brightness coincided closely withthe centre of the optical disk, although the position of the Sun in the sky duringthe observation was the same. Hence the displacement observed on April 30 seemsnot to have been associated with refraction or scattering in the solar corona.
As one can see both the displacement of the maximum radio brightness from thecentre of the optical disk and the intensity of radio emission were significantlyhigher on April 30 (Figure 1) than on April 29 (Figure 2a). On April 30there was also a significant variation of brightness distribution, between 0825 and1025 UT. Such fast variations may be due to abrupt increases of activity at
decametre wavelengths from one or more active regions. This would be in accordancewith an observed growth of solar activity at both optical and high-frequency radiowavelengths (Solar Geophysical Data, 1976a,b). This increase in activity washowever observed in McMath region 14179, and if the same region gave increasedemission at decametre wavelengths one should have expected the shift of emission tobe in the opposite direction to that observed.
98
40 wv (a)440
#44M
3004. 76(b
1440
d 49
/48P
b 4s 4~ 4 -4
Figure 1. Brightness distribution over the quiet Sun, as observed with the radio
telescope LrTR-2 at 25 MHz on 1976 April 30.
99
(a)
152.
/400
f 07 f-9- ~
Figure 2. Brightness distribution as in Figure 1, but for the day of the eclipse,1976 April 29. The circle with centre M shows the position of the Moon.(a) 08h4 5m; (b) 0gh4 5m.
100
2gL7~z(C)
42
/I/
4f 0 3 2 -&7f -2 -
04
14,00
14;6m
.3o% o2 ? f -'
Figure 2. Brightness distribution as in Figure 1, but for the day of the ecl ipse,
976 Aprilj 29. Trhe c irce with cent re M shows the posit ion of the Moon.(c 1 0 1 (d) 11105m.
101
We Concl ude that tile S-component at decametre wavelengths is signi ficant andresul ts in disturbance to thle form of the radio corona, Thus anly estimate of thleequatorial and polar dimensions of the radio corona should take into account radioemission from local sources. It is possible however that thle intensity of deca-metric emission from local sources changes very slowly, as is thle case at centi-metre wavelengths. If so the data above are atypical and are instead related tothe fast growth of solar activity.
Figure 2 shows radio nmaps of the Sun obtained during the part ial silar cli pse1976 April 29. Fihe centres (if the Noon (M) and of the Sun (0) and thtei r disks are
shown. The centres of solar activity are also shown.
Figure 2a is a pre-ecl ipse map, taken at a time when thel Noon was too farfrom the Sun to affect. thle apparent brightness distribution.
Figure 2b refers to the first contact. Here thle Moon screensI. a part ot thlesolar corona and titus causes a displacement of tite centre of the brightnessdi st ribuition.
Figure 2c shows a more advanced stage ol the eclipse. %ow not only thek Coronahut alIso part of the solar disk, inclIutding part of the at t ivy region McMa th 1 4179,is occutlted. Against the haekgrottnd Of the solar corona, the Moon appears as asource in absorption (isolines with "negative" apparent tluxes). Ihis absorptionregion is distinct ly displ aced from the opt ical position of the Moon. Thle ahsorp-tion image seems to he shi fted away from the region of maximum gradient of radiobrightness (ditfference of thle solar corona and lunar brightness).
Finally, Fi gure 2d shows the btri ghtness di siribtit ion at thle maxZIimumII phase Ofthe occultation ((I.66). At this moment tilie Moon 11as cove red both act i e regioenson the solar disk. Thiis has cattsed a marked drop in apparent biri ghitness.
2. METRIC OBSERVA~TIONS
I'he Citlgoora Iratloheliograpit (WilId, 19h7; SheridlAn et tI.19 ) -ktlA%maps the Stil at 80 and 1611 M1z. [act map is derived I rem oliscrvat ions, iott egmut dover about I iin. For a source itt the zenith tite heitrwidth is 1'. al 8f, M11'. and
1.9 at 100 MHz. However, in mapping extended SOtirk LS sit](h JS ih (1, 1 iet SWI
conftusion f rom s itelohes causes thle major loss of dcl mlit ion.
Maps taken cl ose to lotalI neon for Apri 1 29 indl hI it Shown in Fi 4,ures and
4 resp'C t i VtlIy. I'lie ma itn fea ttt r seen a t 80) MHz i s an. *l'igti tion t, t t ,, T 'toutrsa long a I lne W IL i 1i aboitt 1S' o t the equiat orialI plIane of theI) S (it . We be I love,,
(Shteri dan, 1978) that t Iiis sltape, port rays the shape ofit the tona dutring stmspttminimtum. [Ite polar flat tening reflects tite polar corona I Iolts as seeit, dutringsittspet minimttm, in thle white-I iglit corona', likewise the equatorial hbtlge reflectste eqitat orialI streamers. In adit ion we see a sttgges tion of finer st ruittre - two
maxima separated by a shlIlow valleby - on tile disk. Sutch miner maxima art commonanti prohahibly indicate st reamers (hit k and Sheri dan,* 1974).
At 160 MlHz the radio corona is similar in shape bitt smaller than at 80t M117(Wild, 1971). This is bt'cattse most of the thermal brems.-trahhttng is emitted fromlist ahove tile pl asma level , and this is; of course closer to thle phtotosphere at
160 titan at 80 Nilz. As is usual (Sheridan and Smerd, 1977), disk structire is morepronounced at 160 titan at 80 Miz.
102
C. !
I I I I i I
Figure 3. Brightness distribution over the Sun, asobserved at Culgoora at 80 and 160 MHz on 1976 April29. (a) Highest contour 7.8 x 105 K; contour inter-vals 8.6 x 104 K. (b) Hi hest contour 2.2 x 10 K;contour intervals 2.4 x 104 K.
These 80 and 160 MHz maps show little displacement of the radio from theoptical disk. Thus, from this Culgoora evidence, we cannot exclude the possibilitythat the displacement seen at 25 MHz arose through ionospheric refraction.
103
-)' A"('IZ
Figure 4. Brighitne'ss distrihbution over the Sun, asobserved at Culgoora at 80 and 160 MHz on 1976 April 30.(a) Highest contour 6.9 x 105 K; contour intervals7.6 x 10 K. (b) Highest contour 2.0 x 105 K;contour intervals 2.2 x 1()4 K.
104
References
Aubier, M., Leblanc, Y., and Boischot, A.: 1971, "Observations of the quiet sun atdecameter wavelengths - effects of scattering on the brightness distrbution",
Astron. Astrophys., 12, 435.
Bazelyan, L. L., Braude, S. Ya., and Men', A. V.: 1970, "Scattering of the decameter
radio emission of the crab nebula by the solar corona", Astron. Zh., 47, 188,(Soviet Astron., 14, 153).
Rraude, S. Ya., Bruk, Yu. M., Mel'janovskiy, P. A., Men', A. V., Sodin, L. D., andSharykin, N. K.: 1971, IRE Acad. Sci., Ukr.S.S.R. (Kharkov), Preprint No. 7.
Cronyn, W. '. and Erickson, W. C.: 1965, "The flux and brightness of the sun at 26.3
mc/sec.", Astron. 1., 70, 672.
Dulk, G. A. and Sheridan, K. V.: 1974, "The structure of the middle corona fron
observations at 80 and 160 M1hz", Solar Phys., 36, 191.
Sheridan, K. V.: L978, "Some recent explorations of the solar corona fron theCulgoora solar radio observatory", Proc. Astron. Soc. Aust., 3, 185.
Sheridan, K. V., Labrum, N. R., and Payten, W. J.: 1973, "Three-frequency operationof the Culgoora radioheliograph", Proc. Inst. Elec. Electron. Engrs., 61, 1312.
Sheridan, K. V. and Smerd, S. F.: 1977, "Radio and euv studies of coronal holes",
Izv. Vyssh. Uchebn. Zaved. Radiofiz., XX, 1331, (Soviet Radiophysics andQuantum Electronics", 20, 917).
Solar Geophysical Data: 1976a, U. S. Department of Commerce, Boulder, Colorado,
No. 382.
Solar Geophysical Data: 1976b, U. S. Department of Commerce, Boulder, Colorado,No. 387.
Wild, J. P.: 1967, "The Culgoora radioheliograph", Proc. Inst. Radio Electron. Eny.Aust., 28, 277.
Wild, J. P.: 1973, "The Sun - close up of a star", in H. Messel and S. T. Butler
(eds.), Focus on the Stars, Shakespeare Head Press, Sydney, p. 73.
105
01
10. The Rise, Decline, and PossibleRelationship Between Two
Proton Flare Producing Regions£:z Including Details of the Activity
Patterns of Their KnownAntecedent and Descendant
Plages
Susan McKenna-LawlorPhysics DepartmentSt. Patrick's College
Maynooth, Co. KildareIreland
Abstract
An account of solar circumstances (X-ray, optical, radio and magnetic) cha-racterizing the disk transit proton-flare-producing regions llcath 15031 and15266 is presented, together with corresponding details concerning their identi-
fiable antecedent and descendant plages. Differences and similarities betweenthese proton producing centers are discussed, and they are individually suggested
to have constituted responses to instabilities associated with the presencebeneath the photosphere of an anomalous region capable of causing strong magneticflux to appear episodically within the confines of an identifiable solar zone.
107
1. INTRODUCTION
It is the purpose of the present paper to describe in detail the solar historyof two plages which gave rise respectively to (a) the major GLE proton flare ofNovember 22, 1977, relevant to STIP Interval IV and (b) the well known sequence ofoutstanding proton producing flares occurring in April-May, 1978. In the formercase, the energetic flare occurred during the first disk transit of an active centerthat first developed on the invisible hemisphere. By its next appearance at theeast limb, this region was of negligible solar importance. For details see Section3. Plage 15266, progenitor of the April-May sequence was, on the other hand, asso-ciated with the emergence of new flux within and close to an active area that hadalready shown as many as five transits and the magnetically metamorphosed regionthus produced thereafter went on to transit a total of five times. Section 4 pre-sents not only an account of solar circumstances attending the transit of Plage15266, but also those of its various known antecendents and descendants. In Section5, differences and similarities between the two proton producing centers are des-cribed, and in Section 6 they are individually suggested to have constituted respon-ses to instabilities associated with the presence beneath the photosphere of ananomalous region, capable of causing strong magnetic flux to appear episodicallywithin the confines of an identifiable solar zone.
It may be noted that "returns" to the visible disk of individual plages couldbe recognized with reasonable confidence by giving due cognizance (a) to the effectsof differential solar rotation and (b) to evolutionary changes associated with thegrowth and decay of underlying spot field deduced, from white light records, to havetaken place while individual regions were transiting the invisible hemisphere.
In order to summarize the waxing and waning of activity within transition plagesrotation by rotation, an "activity profile" was prepared for each transit describedwhich shows, In tabular form, the day by day evolution, as monitored in calciumlight, of the plage concerned, as well as information concerning variations in thefield strength, area, spot count and magnetic classification of its underlying spotgroup. The frequency, with respect to class, of any observed flares is also givenand the numbers of Sudden Ionospheric Disturbances, X-ray bursts and single frequencyand dynamic radio events attending the flares mentioned are listed.
For convenience, those anomalous locations within the solar atmosphere in whichthe proton flares of November 1977 and April-May 1978 developed will hereafter bereferred to as Active Regions I and II.
2. SOURCES AND KINDS OF SOLAR DATA ANALYZED
The experimental data used were taken primarily from the monthly publicatonSolar Geophysical Data (SGD) compiled by the National Geophysical and Solar-Terrestrial Data Center, Boulder, Colorado, U.S.A. Errata or revisions to primarydata sets which appear in these publications from time to time are incorporatedwithout comment into the material presented.
Flare data were taken from SGD "Comprehensive Reports". Flare importance isgenerally deemed to be that assigned in the associated summary "Group Reports" but,in individual cases where no importance rating was originally ascribed, a suggestedimportance, based on such flare parameters as were available, is here supplied withan added question mark to indicate special uncertainty.
108
Through the kindness of the staff of McHath-Ifulbert Observatory, center line aswell as on and off band Ha pictures of the April-May flares were available forstudy. Data concerning calcium plage regions were taken directly from the daily SGDreports of the McMath-Hulbert and Catania Observatories. Sunspot data used includesthe tit. Wilson and NOAA (Boulder, Ramey, Manila) SGD reports. Original drawings ofthe individual groups made at Mt. Wilson (and supplied by World Data Center A, NOAA,Boulder, Colorado, U.S.A.) were also consulted and used where necessary to correctthe SGD classification listings.
Information concerning the X-ray variability of individual solar regions wastaken from SGD "quick-look" OSO-8 maps. These displays, some of which are repro-duced in Figure 1 were constructed from coverage obtained during a representativethirty minute period (neither the most quiet nor the most active).
CA., SO0 C.fto. led, C A., '602
OCT. 20 NOV, 21 DEC I?1410 U. SIS U.T 5so0 U,
ALMAT"-CJL'T CALCUMt MAT- B7.AT I CALCIUM R T ftlI.ALM-" CACLOM
o./ oo., O
0"0) Or
NOV 20 NOV. 22 ov 2309 , LT 00 DO U 1 30 UT
..0C01CE 050-6 X-RAY LOCIHED No 0 8 X-Re o AY LOCTIED N o SU*0S .- RAI
10
15 5s 30 '5
a
Figure 1. Top: Calciun plage reports of the fcl[ath-Hlulbert Observatory showingthe same general region of the solar atmosphere during Carrington rotations 1660,1661 and 1662. Plage 15031 (cf. lb), which returned to the disk as plage 15063in December, was associated on November 22, 1977 with the occurrence of an importantproton event. Bottom: OSO-8 X-ray maps taken before, on the day of and followingthe proton event. The relative-intensity of a region is here indicated by dots ofdifferent sizes where "detectable", (D), nominally corresponds to 5X times the
background value. The three dot sizes employed represent D-20D, 20D-500D and> 500D. If the source varied by more than an order of magnitude over a time periodof < 2 hours, a "highly variable" indicator, consisting of a dot surrounded by acircle, is used rather than a "typical" intensity.
109
Satellite data concerning X-ray bursts and world-wide reports of SIlD's weretaken from composite SGD lists. Suggestions contained in these compilations as tothe possible association of individual events with flares are altered without rommentill the text when updated information rendered this necessary. Similar correctivetreatment is accorded to SGD lists of "flare-related" radio bursts recorded at fixedand swept frequencies.
Data concerning solar proton events were taken from SGD "provisional" lists.
3. DEVELOPMENTAL HISTORY OF ACTIVE REGION I
Nctive region I was first identified during Carrington Rotation 16h1, whenit was designated McMath Plage 15031. Only adjacent Plage 15033 had previouslycrossed the solar disk, i.e. once as Plage 15015 (see Figure 1), and this latterfeature was already in an advanced state of decline and showed no special acti-vity during the November transit. Structures 15028, 15034, 15036, 15037, and15040 had, on the other hand, all developed since the previous rotation so thatPlage 15031, which first rotated over the east limb on November 12, 1977, deve-loped at a previously "quiet" location during a period characterized by ambientrising flux.
3.1 Plage 15031
OSO-8 observations indicate that Plage 15031 constituted a weak X-ray sourcetip until November 18, by which time it displayed a moderate intensity enhancement.After an ensuing decline it became "highly variable" on November 22 (the day ofthe proton flare), returning to a steady condition by the succeeding day. (Com-pare maps (e) and (f) of Figure 1.) These events were complimented by a risinglevel of sub-flaring in the plage previous to November 22, an effect apparent,(see Table 1) despite significant gaps in the flare patrol record over the daysconcerned. This waxing solar activity culminated on November 22 in the produc-tion at 0945 UT, of an importance 2B flare accompanied by an importance 2 Sudden
Ionospheric Disturbance (SID) of Widespread Index 5 and broad band single as wellas swept frequency radiation, (see Table 2). An associated X-ray burst was alsorecorded. These related occurrences were climaxed, in less than an hour, bythe onset of a ground-level particle event.
It is noted that three of the flares occurring on November 22 prior to theimportance 2B event, were associated with the generation of minor centimeterburst radiation as detailed in Table 2. The closest recorded event to the majorflare was additionally associated with the production of an importance 1 SuddenPhase Anomaly (SPA) of Widespread Index 1. The first flare to be associated withan ionospheric disturbance previous to this was an importance sf event at 1132 UTon November 20 (accompanied by an importance I Sudden Enhancement of Atmospherics(SEA) of Widespread Index 1). After the major flare of November 22 no furtheractivity other than an importance sf event on November 24 was reported in region15031 prior to its limb transit.
Examination of the original 'It. Wilson drawings reveals that by 1645 UT onNovember 17, a B(6) configuration was present in the underlying spot group.(Corqtare Flgures 2
a and 2b.)
110
it. 'o"4 'OJ
I~~~ 8) 0) . 8 8 .- )) u 0 8) 0 C
m 0 1~ z-8
- 0 8 4 0..'04-41 0C ) .~
.c 40 4 04)
cu " 8 4) C'4 "0 :j4 to8) 80 To WO to. 4- 0
w c OD 4) .0 0 ~ 8>28 cw0 0 $4 $W8) C 80 8 8) 0 -CL 'v 4)o
CC w ~ f 0 1 ab 0)C
Cw C - )-4- 0m=. r F= s.I0
4- -)4' cc)~4 a.40 0.0) 0 m- -
005 4 -8 *..'C.m
) 0.8 .0X0 0 >j a. 8)4-4
a-00) mU
04-0 U W C
3t w~ 4 0)C
80- *a4E ) CA18V) a) 0 -m
8)00 0E 08r
C) 8)8K1 0 w10
.0 w )8C m).. m8
Wm w) 18 0 'It, 8) "
JJAQ)J -n44m
-'0 0.80 a >
2 m- a) >.' a- )8 .
n 'A0C)W)o j 8.-aa, 0 4a ) c8 'a)1)
Q) 0 rm' a 0 CU C) 8
c 8 8 0 )' I--4- 18)D r 0 o
C)) (D 80 0 4~4- ) ) 4- - 8 0. .. lC I 0
0 >.C x 8).4. L. tofl-4
8) 0 0 Ci 4> ') a4)
0 w a 4 0) w 7-1 ) 8) . ..
8) a u) 8)0 w4 '48)'u 8 co -8 8 EM
80 )' 0y -,4 )-,C4 Z5) 8 C ' .
A I
0 00o 3 'n ;g4t
x~I t1 oo't41 - .r4 VI* .j~ 01
u2-w 0
N .0
St? 04.
-P4lo
'A4 0 -nC T m I7'
wo In 0- 0 0-06
a. 1 0-
41
* 64 4441 44~ a ' U N N ~ ~ 4c
U4* ,2dg,, jg jii~o112~
The next available record, made at 1610 UT on November 20, shows that, by this
time, the following spots had become dominant. No magnetic: observations were madeon the day of the major flare. Thereafter, the spot group was seen to be in rapiddecline, reaching class a f by November 24. 1 report from Boulder indicates thatthe largest spot in the group had an asymmetric or complex penumbra on the day before.he major flare and a rudimentary penumbra on the succeeding day. (See Table I fordetails.)
9.o00 UT i 4 U T
20.0V20VII
423
Pp
E30
W ~0
OV. 20 Nov 2116 10 U 1 16 40 UT
R*1 4 147 o s\*S 4.3 .24---
R6
PI)
NOV 23 NOV 241640 UT 95 23 UT
090
.24 --- --- .. 24
V13 ",. ,W
1.) M I
Figure 2. Reproduction of drawings of spot group
Mt. Wilson 19894, originally made at the Mt. Wilson
Observatory during Carrington rotation 1661. Thenorth is at the top of each drawing; the west is tothe right.
113
3.2 Plage 15063
On its next disk transit, when it was re-numbered McMath Plage 15063, the re-gion which had previously been the source of the proton flare was so substantiallyreduced in area and brightness as to be of negligible significance, (see Table 3).No observed flaring occurred during the course of its transit, and the Mt. Wilsonobservers reported no spots in the underlying magnetic field.
4. DEVELOPMENTAL HISTORY OF ACTIVE REGION II
Plage 15266, which gave rise to a spectacular series of proton flares inApril-May 1978, occurred in a part of the solar atmosphere that had been highlydisturbed since the preceding January. As shown in Figure 3, which traces the
Transit 1 .................... 15097 (N20, L62) C.Rot 1663Mt.W. 19931. Class J
Transit 2 ...... 15135 1N20. L731 15134 (N27, L73) ............. Transit 1Mt. W. 19950, Class a Mt W.19%9 Class DMt W 19951. Class at. I
Transit 3 ....... ........................... 1517 2 (N 20, L71). ................... Transit 2 C.Rot 1665NOAA Class EHINOAA Class DSINGAA Class 0SI, - I
Transit 4 ...... 15214 (N 23. L78) 15220 1N31, L59) ............. Transit 3 C RotMt W. 19965. Class a Mt. W 19990, Class X"t W 19994. ClassI I
Transit 5 ............................ 15266 • 15266. N23, L68) .. ....... ................. Transit 1 C Rot 1667Mt W. 20018. Class D (Py)Mt.W 20019, CLass D (p)Pt W 20020, Class O pf)
15314 (N 18, 179) ............ Transit 2 C Rot 1668Mt W 20044, Class DMt W 20048. Class 01Mt W 20049. Class yI.15368 (N 19, L81) ............. Transit 3 C Rot 1669Mt W 20088. Class ocPMt W 20091 Class D(Pj)Mt W 20098 Class airMt W 20100 Class fJ
1154,20 (N 19. 166)........ Transt 4 C Rot 1670Mt W 20142. Class attMt W 20150. Class P3
15473 (N21, L73) ........Transit 5 C Rot 1671MtW -
Figure 3. The development and decline of that active solar centre which, inCarrington rotation 1667, gave rise to the spectacular series of proton flares ofApril-M"ay, 1978.
114
-. , ... 2'o r ca
,- ,-, 0 •.- w, .C u w r
V u 0 c w Ecc0 Q) W4) 4).
W.0 0 0 wAcr, W J~ .
-4 Q) eC/ - o)~ 0
00 u t o
a to m -V CL m ) )
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0C. 0 0
0
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r- r I0. J A tc o T 0.o W
> UW CC-
M 0C 0 0 C
Q) " m" w ' C '> to w''
-. 0. . , W C O. . .... ...
0.. C- C > a .
00 0,,0, "-
U)C
0 J 0i -11 1)( a> C: -w fl
U W O * U . , C , 0 . . . .. . . . . ... . ...
M- r. 'A
r 'r
It'-W -'--CCC . ~ -~115
development and decline of this anomalous region, the antecedents of this structureincluded a sequence of four reappearances of Plage 15097. Another sequence of threereappearances of Plage 15134 occurred so spatially close to the first set mentionedthat, in Carrington Rotation 1665, contributions from these two sources were jointlygiven the number Plage 15172. During the next transit, the individual elements ofthe two adjacent plage streams were again separately distinguished.
By Carrington Rotation 1667, the underlying magnetic situation had again becomehighly complex. New flux, which had in particular emerged within the ambience ofpreviously existing magnetic fields, had indeed so fundamentally altered the char-acter of the overlying plage that, while vestiges of its previous structure (desig-nated in Figure 3 by the number 15266') survived, the overall character of the regionwas so modified that it was deemed appropriate to consider it to constitute a "new"feature (15266 + 15266'). This "new" plage made four further appearances beforesubsiding into the ambient background.
4.1 Plage 15097
The first member plage of that series preceding the April-May transit rotatedover the east limb on January 7, 1978 during Carrington Rotation 1663, when it wasaccorded the designation Plage 15097. Its passage was characterized on days 11-14by a moderate level of sub-flaring (probably incompletely reported due to extensivegaps in the H-alpha patrol observations), falling off thereafter to cease on January18, (see Table 4). Associated electromagnetic accompaniments were minor. A cm-draburst of maximum measured mean flux density 9.2 solar flux units (sfu) at 26, ,ilz,accompanied an importance sn flare at - ?240 UT on January 11. A low level o) X-rayactivity was reported by OSO-8 observers in the plage on the morning of January 11,declining on suceeding days. lowever, a "highly variab " X-ray condition was re-ported to be present close to the time of an importance '-?) fl,t .i the regionat 1005 UT on January 16. Associated Mt. Wilson gr" , 31 rose in (_mplexityfrom NOAA Class BXO on January 9 to Class DSI by at .. it Jinuary i3.
.t.2 Plage 15135
Plage 15135 constituted, on February 1, the first return on Plage 15097. Thegeneral enhancement in area and intensity attained during this passage was a con-sequence of the emergence at this location of new magnetic flux since Mt. Wilsongroup 19931 of Class AXX which transited the west limb on January 18-19 had by thistime re-emerged to the east as developing group Mt. Wilson 19950, accompanied bydeveloping group Mt. Wilson 19951, (see Table 5). OSO-8 observations indicatethat, from late on February 2 through February 3, the plage was "highly variable" inX-rays, declining on the succeeding days to a steady but moderate level of activity.February 3 was also the most flare active day of the transit. Three of the minorflares then occurring were collectivly associated with the production of 3 SID's, 2X-ray bursts and 2 cr-di bursts. These latter events had durations of only 0.4minute and 4.8 minutes with corresponding mean peak flux densities of 10.1 sfu and3.3 sfu, at 2695 and 2800 Mliz respectively. These flares occurred during an intervalwhen the underlying group was waxing from NOAA Class HSX to Class CiO.
116
m~ w 0r).. - -wC0
,v C-' m a wT~ ~~ 0. . )au- 1. - o-r u -
w ) L 1 a) w. a
a)'5 a W00a)w 2)
1 - 02 '
41-a M )
r-- r) o))'aa
- ' ) I' - ) a 0.
a) r aa )
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Z22)) .')10 : ~ ~ ? "'a : ?'w I N.
a)( -K
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".Ca Q L It,0
'5*-0 a) .->)2 CaL.-0
- u 0 c M~a - "- -
C. 1- ;-5 0 )C) C) 'a.-~ ' 0'o)5.'0 a1 w '~
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cu to u
a)4 r to 4)
0(W c -. r .
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a.' r~
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ODL
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Ent
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00 4
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4-. w0 IV-
44118
Thereafter, with the exception of an X-ray burst and an importance I SPA ofWidespread Index I recorded on February 5 in association with the occurrence at0014 UT of an importance sf flare, no other event recorded during the transit dis-played associated activity.
t.3 Plage 15134
New adjacent Plage 15134, which came around the east limb on February 1, wasassociated with developing hi-polar group Mt. Wilson 19949. From approximatelymid-transit, its characteristic of constituting a "highlv variable" X-ray source wasconplinented by the corresponding close association between rather minor flaring inthe plage and the production of time associated X-ray bursts and SID's, (see Table6). Sometimes also, centimeter or centimeter-decimeter bursts of low mean peak fluxdensity occurred in association with the observed flaring.
Examination of the origiaal !t. Wilson drawings reveals that, on February 8,ahen Plage 15134 was showing particular variability in X-rays and the region wasespecially flare active, a (6) configuration was present in the underlying spot,,roup. No later spot drawings of the group are available, but NOAA reported a de-cline from Class ESO on February 8 to Class DSO by February 13.
An importance I+ SID of Widespread Index 5, accompanied by an X-ray burst,broad band cm-dm radiation and Type II and Type IV events, which began on February10 at 2103 UT during a gap in the optical observations, may also have been asso-ciated with unobserved flaring in this region-particularly since 050-8 observersrecorded a very dramatic peak in the level of X-radiation produced by the plage at0100 UT on February 11. No other active center on the disk showed a correspondin;level of X-ray activity.
t.t Plage 15172
Plages 15134 and 15135 were deemed on their return during Carrington Rotation1065, to form a single composite structure designated Mc}iath Plage 15172, (seeTable 7). Observers at Boulder recognized three spot groups as associated with thisregion at approximate mean latitudes N14
°, N20', and N27'. The level of X-ray
activity as observed by OSO-8, as well as the frequency of X-ray burst and SIDassociated flaring within particular parts of the plage, appeared as could be deducedfrom those magnetic reports available from Harch 5 to depend intimately on thelevel of complexity of the immediately underlying spot group.
The most important flare to occur during this transit, a 213 event on March 6 at1123 UT accompanied by an X-ray burst, an SID of importance 2 with Widespread Index5, a broad hand radio burst with mean peak flux density 50 sfu at 2650 'IlIz and aType H1 event, occurred above that group located at N27
0 during an interval when it
was waxing from NOAA Class HSX to Class CSO. Although this latter group had declinedto Class HAX by March 9, it revived to reach Class DSI by March 10. The most en-ergetic overlying flare to occur during this interval was an importance sn event atD)823 UIT on March 9, accompanied by a broad band radio burst of mean peak flux density4 sfu and duration 101 minutes at 3100 Mliz.
119
-I M )0 C )-
c4 -. 0.4m)00r-00 m r) m)0 4
CL. 4) 000 -a, )~4
w 00 I..O
v u )a o .
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0 4 . )* 4 0 2 4 . .. . .. . . .. . . .. . . . . . .
0- . I.,,i. 2.
0'-' ~ c 0. 4 00 4 ->0 r w ) 0 _ 0.,-0 -.. . . -2 L. .4) 0 W
4)00 4).-- 00". i . .. . . .
0 r-. r ))) *04
200 0 C)0 0 4 a) a) 1 4) 0 4
.. . . ..- 4. . . .. .0. . ...0l l - i r .i &"-.
mw 1100,C..U° M Q;
,)04,4,4)),. ,*, - ,
4) o 04 'o~' . .. .. ..
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z 0. 4). ) 0..
0. 2 - 04 4 , 4) ,, , ,, - -, -'C,
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000 9= ) 0 ~ - ''' ' '
lu
I-' 40 2 04) ;
0 w0
.. . ..0..-. >.-.0 4 >4 -4
0.-4 m)'Q. 0 wO
4) ~ a 00o m0) .. 4
*.-.) 4a ..- >
r I-. 4)00
4 -4hJ~ .. 4,).4.. z)
a) LI) 4. 04 w4) I03In w)4)1. m)..
U) m '0-44)) w ~4) .0)4 '- u-1- C 4) 4
m0 e 0)04c w-- 0t q)C o 0>4) 0 c- m0 4' x )
00 a -) 0 - 4) u
>C -j C*
- u0020 120
TABLE 7.
The evolution, as monitored in calcium light, of those plages in which thle proton
flares of November 1977 and April-May 1978 developed (cf. Tables 1 and 10, re-
spectively) together with corresponding details concerning their antecedent and
descendent plages. The sources of these data together with complimentary information
concerning plage related magnetic fesLures, flare activity, S10's, X-ray bursts and
single frequency and dynamic radio events, are described in Section 2. The conven-
tions of presentation adopted are generally similar to those contained in Solar
Geophysical Data (Supplement, No. 390, 1977. The number of flares identified within
a particular plage, however, having an importance class with the range sf to 4b *s
indicated according to the scheme (number of flares identified (reported flare
importance)1. In cases where no importance rating was assigned in the source
material to a particular flare, a suggested importance, based on such parameters
aos are available, is assigned by the author and the event listed according to
the scheme I? (i.e. suggested flare importance). The areas are measured in
millionths of the solar hemisphere.
151'2 Ft.*G t (( -0 3004 00 I_ 0 l000 4000: 100 7N00 .000: @70400 500: 3200 1100:
433361' . .0 0.5 1., 3.5 1 .5 0. 1.53 5 .5 .3. I ., 3.0 2.0 3.0
.20, 40 3(0t. 080 320 J"36 200 . 340 30 340 0
~~0-- -- -03-3 . - -2. 2....2.3
or'C.o3411 35 04
_E.2' -40.0 :50 60 30 .30 330 50
0(40.. . ~ W33: 2:33 :30
63 I LASS . 6 06 :10: 6 .330:33:00 00
Slot 104 4 30 43 2640.'
--- -- -- -- 4..3 ...6331. . . .... . .
00. 46 . S645W ICS A 1D.D.o .c --:-- -- --- :- --
-- --- --- --- - - -- --- - ---- --- -- -- - - - -- --- - --- -- -- - --- --2- -- --
4.5 Plage 15214
The inherent twin elements of composite Plage 15172 of the previous passage,were assigned numbers 15214 (rotation 4) and 15220 (rotation 3) during their 'arch-April transit, (see Tables 8 and 9). Plage 15214, which had shown some X-rayvariability during its previous passage, was again seen from OSO-8 to be variableduring Carrington Rotation 1666. Observations were somewhat limited at both X-rayand optical wavelengths but, at least from April 7-9, a fairly high level of minorflaring, accompanied in particular on April 9 by X-ray burst and SID activity, wasrecorded in association with this region. A centimeter wave length burst of duration4.3 minutes and mean flux density 8.1 sfu at 4995 MHz, was additionally associatedwith an importance sn flare at 0215 UT on April 9.
Underlying bipolar groups Mt. Wilson 19985 and 19994 appeared, from the avail-able measurements, to have been of low field strength. Also, the plage was in astate of general decline as regards its area, and probably also its brightness, asit transited the west limb.
4.6 Plage 15220
Plage 15220 was in rapid decline as regards area and brightness during itsMarch-April disk passage, (see Table 9). It produced several sub-flares while itwas close to the east limb but no accompanying electromagnetic effects were reported.
4.7 Plage 15266
It night reasonably have been expected when Plages 15214 and 15220 transitedthe west limb during Carrington rotation 1666, that the overall trend of decay ex-hibited at these locations would be maintained. However, the emergence while thedisturbed region was on tile invisible hemisphere of new magnetic flux, particularywithin the ambience of previously existing fields, so radically transformed tilecharacter of those plage elements present that, as already noted, the resultantstructure may he deemed to have constituted a new feature (15266, incorporating butvestiges (15266') of previously recognizable plage components. See the overall"Activity Profile" presented in Table 10. Spot Groups Mt. Wilson 20018 and 20019transiting beneath the "new" plage appeared to .: returns of Group 19985 and 19994of th, previous rotation, although the birth of new flux had imparted to thesepreviously decayed fields a greatly enhanced complexity. 'It. Wilson 20020 on theother hand was located in a previously undisturbed solar region to the east of theother groups.
Examination of center line and off band H-alpha pictures of the fi.ous seriesof proton associated flares occurring in the plage on days April 28, 29, 30. and MtayI respectively, reveals that the events of April 28 and 30 avoided the large leadingspots of 'it. Wilson 20018 and showed a preferred association with Group 20019, whichdisplaved a 6(8y) configuration on April 28. The flares of April 29 and May 1 onthe other hand were associated with (6) Mt. Wilson Group 20018. See the sunmary of
122
- C
<' c A
-r r404
C-0 w
0' 0 0
4 . 0 ,. -
w t.4 44 4
.0 w .
00--0 v
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0 00.
0 -4.'0 ,
ooo
.2 c :, .i 4 , 4'0 0) 3 .- , ,.123
mC C) 0 0. 'o a,
4- 0 IUD U - )4
0. I) C ) '.' U 44
't 1 1 ..0 .. ..ow0 =4CCa)oZ00010(4C Cx m- 4.0) w
0 ) c tz Q) C 0 -V 4)
0..' w- uuCU-. a) ) CA M' M- C) C).
0C4 U '- (0 .v.40C) CO a. 44 - SCO4
C) U- C ) .C C 0 ) C
-0 0 '0 cd C w).C)%
C)O 0 0- ..
C) Q ) 0 CO 0 0 C
0. C -. C c -4..
OC.4 C It 400C
.UC 4- >00(4.
O~cm- 0 Lo.
0041 > ~ 44'C 04
0,I,, )W r0C -C 00a
C 4 0 ).4-..aj) C.44~)). . . . . . I.
>. -0 zo )> Co OOU U WO)0.,
. .'! v w . . ) C)- V.. C .C ~ 4.
C)C CC) ."!124
TABLE 10.
The evolution, as monitored in calcium light, of those plages in which the protonflares of November 1977 and April-May 1978 developed (cf. Tables I and 10, re-spectively) together with corresponding details concerning their antecedent anddescendent plages. The sources of these data together with complimentary informationconcerning plage related magnetic features, flare activity, SID's, X-ray bursts andsingle frequency and dynamic radio events, are described in Section 2. The conven-tions of presentation adopted are generally similar to those contained in SolarGeophysical Data (Supplement, No. 390, 1977. The number of flares identified within'a particular plage, however, having an importance class with the range sf to 4b isindicated according to the scheme (number of flares identified (reported flareimportance)}. In cases where no importance rating was assigned in the sourcematerial to a particular flare, a suggested importance, based on such parameters asare available, is assigned by the author and the event listed according to the scheme1? (i.e. suggested flare importance). The areas are measured in millionths of the
solar hemisphere.
3 3 , ,, 0 f 3
.... 3.. u. ,' 2: 25 2t -3 .24 23 :30 . 2 . 4 5 . 3 3 : 8
1-00 .. , 550 W , 1. - 1100:3500 12100.323000; 3O0:1300-2330.3 : .00 500. 2000
.5 .5 1.5 3.5 3.5.4.5 3.5: 3.5 3. 3.5 3.5 35 35 .
-4,[ 33. , 8 .0 , 43. S. 430 .4l3I0 f 1150 740 a" 1080
Eo -,0i " 1 1
- ~ ~ ~ ~ - :..3 .' 4 4:3
-, , ,L, p, , I .. , 3. s , , 3l , .5, 31 j, . 4 ,
3.. .. f.4/3. sl. , 4 43.. .f. 3., 3(43• 33 3,3. 5.o- i.3
343'. l.s ] i i 3.3 433. 3333, '4 '3. 3.-3.-3
.u~c.-. .3,3. 3: 3'3: 3 3383 ] I
1 2 1
40. P4-4214
'3343.~ 2.-3 1~3 1,3
-- -- -- -- -- -- -- - ---- -- -- --3- -3- ----------- --------- -- -3
125
information concerning these various events contained in Table It.*
An importance sb flare at 0010 UT on April 28, accompanied by an importance I
SID of Widespread Index 5, broad band radio radiation with maximum mean flux dnsity'167 sfu at 2695 Mlhz and a Type IV event, also appears to have been associated withGroup 20019.
On 'lay 1, a broad band radio burst with maximum mean flux density 3.4 sfu at
2695 hz and an accompanying Type II burst, occurred in association with an import-
ance sn flare at 0538 UT above hIt. Wilson Group 20020. This latter group was waxingfrom Class 8 p to 6(ay) configuration on this day and, thereafter, until Plage 15266transited the west limb on May 9. All of the energetic flaring associated with thisregion (see the summary list in Table 10) occurred directly above the aforementionedgroup. OSO-8 observations indicate that the overlying active part of the plage,while constituting a variable X-ray source on the days previous to transit, becameoutstandingly variable by at least 0315 UT on 'lay 7, that is prior to an importanceIn flare at 0327 UT accompanied by an importance 3 SID, and X-ray burst, broad handradio radiation with maximum mean flux density 250 sfu at 2695 hz and Type 11 andType IV events.
4.8 Plage 15314
Plage 15314 was a return of Plage 15266. During this transit, in CarringtonRotation 1668, it was associated with underlying Ht. Wilson spot groups 20044, 20048,and 20049, (see Table 12). Of the groups concerned, Mt. Wilson 20048 attaineda 3y configuration on May 27, declining thereafter in complexity to a bipolar state.Group 20049 showed more variability, first attaining a y configuration on 'Iay 27and after a general decline reverting again briefly to Class Y between May 31 andJune 1. These variations were complimented by corresponding changes in the levelof (localized) overlying flaring accompanied by X-ray bursts and/or SIlls, and in theability of part of the sipernatent plage to function as a "highly variable" sour-eof X-rays. In some cases minor radio bursts accompanied the flaring, and in particu-lar on lay 28, a broad band burst with mean maximum flux density 5 sfu at 2695 ,'llzaccompanied an importance lb event at 1310 UT.
'It. W1 Lson 20044 had attained a delta configuration by at least 'lay 2S andretained a high level of complexity until June 1. Again, special periods of X-rayvariability in the overlying plage, as detected by OSO-8, were complimented by cor-
responding variations in the level of X-ray burst and SID associated flaring producedat this location. Broad band radio bursts of low peak flux density additionallyaccompanied "ionizing" flares at 0630 LIT and 1523 LIT on Mlay 30. The most energ;eticevent to occur during this transit was an importance 3b flare at 1006 UT on Ia',, 31,accompanied by an SID of importance Class 2+ and Widespread Index 5, and an X-rayburst. Its accompanying broad band single frequency radiation achieved a mean ax-inum flux density of 49.2 sfu at 2695 Mliz and associated Type I and Type IV bursts4ere produced.
Examination of underlying Group 20044 reveals that, between May 31 and Tune 1,its internal magnetic gradient decreased significantly. A further importance lbflare above this location at 2113 UT on June 1, was reported in addition to ionizingand very minor broad band radio radiation, to be accompanied by a Type II burst. By
*A forthcoming paper (in preparation) will suggest that these were not in fact
four independent flares but rather two complex flares each with two phases andtwo maxima.
126
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A I
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~
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127
TABLE 12.
The evolution, as monitored in calcium light, of those plages in which the protonflares of November 1977 and April-May 1978 developed (cf. Tables 1 and 10, re-spectively) together with corresponding details concerning their antecedent anddescendent plages. The sources of these data together with complimentary informationconcerning plage related magnetic features, flare activity, SID's, X-ray bursts andsingle frequency and dynamic radio events, are described in Section 2. The conven-tions of presentation adopted are generally similar to those contained in SolarGeophysical Data (Supplement, No. 390, 1977. The number of flares identified withina particular plage, however, having an importance class with the range sf to 4b isindicated according to the scheme (number of flares identified (reported flareimportance)). In cases where no importance rating was assigned in the sourcematerial to a particular flare, a suggested importance, based on such parameters asare available, is assigned by the author and the event listed according to the scheme1? (i.e. suggested flare importance). The areas are measured in millionths of thesolar hemisphere.
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128
June 2, Grou1p 21004-,3 1.i1 l i ivd t 'Cigiet i- (:1 a/ind nii further eneorget ir
flares ujer, seen ti he si-ji with1 this re,;ion. A Forhiioh dcrease of ')/ whiclioc-CUrred near nd d-day in It ine M, ml V)t hIave hcen assoc iate d w Ith thle act ivity les-c r ied in thle prev iii)s iir.izr tpls.
4.9 Ptdge 15368
P lage 1 3568 w~as a retuirn of P lage 1331.3 in Carr ington Rotation 16)69 anti wasassociated during this traverse with 'It. !hilson Group 20088, 20091, 20U98, and2010)0, (see Table 13). Of these, the bipolar Group 20100 marked the position ofemergence of new flux wile Groups 20091, 20(188, and 20098 corresponded with pre-viousiv transit ing Groups 20049, 20044, and 2018.
The f lire patinol was very inc omplete durini; nuch of the relevant per iod, part i-ii tar l onl June 16 and from June 22-27 incluisivye. On Junet 22 the largest recordedflire of the transit, an imortance 26b event at 1 o-3 UiT occiurrel., accompanied by animpor tancP 2+ SID of ',i despreid Index 5, broad hand rad in radia1t ion with naximum
flux1ii ilensltV 31.0 sfuj at -'f,9 nli and associated dyniamic e~vets of Type 11 andIV. s11) relaited f lar ing, wi th some acompanving miinor radio hurst ic tivi tv was alsodetected On June 20.
Fximiqtat ion w the 'It. *jilson drawings reveals thiat the imnediatelv underlv li-Ls nit proup 'It. i soi 2'W91, achiieved a 6 By 'con fi1gura t on he tween Jun(, 19 and 2UL,decliiing toi a i cini oni by .June 22. There wau a1 revival in anei complex-it' hvtweken Jine 2 )and 2ot wheni a y configiirit ion was aIaLT-nattained. The availlhle'records indicite the conplimeintary presence of mioderato flare aictivity, aconpaiiiedhy low, ener*' ioicin), and/or radio radiation, from loine 27 to .11,I0e 3,il .\ISO,oS-obhser vat ions Show that the re levant part of the active '1 aeT whlto sos Li (ned lvcons t itultinrg a n X-ray source dIir in g i ts1 June)t tr a ns it , b, ,c a e -I igh 1Nv var i i blIe-bvtwee, at least I 03: r7 On -Nune 28 and 01 15 FT on Jujne 29. -Ir un-der I'v in)! spot nroupagain decayed tro)m cent guraitinn By to conf iguration 3 1) hetween Junie 23 and 29.
LIO 0 Page 1.-)420
P'lige,, 1342(0 was a return of P la)ge 151SJ1i in Cirr ingtoni ROtat inn 11 7i> It wshowever, sii reduced in hr ightness and in area is ti hi' a mi nor solar feature, (see,Ta-ble I4) while associated Mit. Wilson Griiips 21)12 and 2110 ere hut vet~iges ofthe corrvspondi n , main groups of tihe previous passage. 001.. one faint sub-flare?reportedt In the region diiring this transit anid no :issoc iatid N-rayv act ivi tv wasreo irded.
HI1 Plage 15t731
Place 15473 was a scarcely perceptible return if ('lage 15420 duritn, carringt inhlitition 1671, (see Table 1)). No flare of Magnetic icti -itv was reported ii: thisreOgio~n.
129
TABLE 13.
The evolution, as monitored in calcium light, of those plages in which the protonflares of November 1977 and April-May 1978 developed (cf. Tables 1 and 10, re-spectively) together with corresponding details concerning their antecedent anddescendent plages. The sources of these data together with complimentary informationconcerning plage related magnetic features, flare activity, SID's, X-ray bursts andsingle frequency and dynamic radio events, are described in Section 2. The conven-tions of presentation adopted are generally similar to those contained in SolarGeophysical Data (Supplement, No. 390, 1977. The number of flares identified withina particular plage, however, having an importance class with the range sf to 4b isindicated according to the scheme {number of flares identified (reported flareimportance)}. In cases where no importance rating was assigned in the sourcematerial to a particular flare, a suggested importance, based on such parameters asare available, is assigned by the author and the event listed according to the scheme1? (i.e. suggested flare importance). The areas are measured in millionths of thesolar hemisphere.
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5. IIFFERENCES AND SIMILARITIES BETWEEN TIlE PROTON PRODUCING REGIONS
Plage 15031, parent to the November 22 proton event, which was born in a pre-viously "quiet" solar region, remained identifiable during only two solar rotationsand constituted a feature of negligible solar importance during its second traverse.Plage 15266 on the other hand, parent to the April-May series, was born among theremanents of a previously existing active region and continued to retain its newidentity over five full rotations.
In both instances, an anomalous magnetic situation characterized that part oftie photosphere directly underlying the active plage. No white light pictures wereavailable on November 22 but delta configuration and reversed polarities were re-ported to be present in associated Mt. Wilson Group 19894 on preceeding days. Inthe case of the April-May events, the birth of new flux inside the boundaries of analready existing, but greatly decayed, magnetic region, resulted in the build up ofan especially complex magnetic situation, and associated Mt. Wilson Groups 20018 and20019, both displayed a delta configuration by (at least) April 28, (the day of thefirst major proton flare).
An impression is given by the records that the occurrence of the November 22flare may have been associated with an accelerated decay of underlying Group 'It.Wilson 19894 which had declined to an a f configuration by November 24. In the caseof Groups 20018 and 20019, the intermittent emergence, within and close to the same
general location, of further flux over a period encompassing several subsequentrotations, counteracted any incipient general tendency towards magnetic declineafter the major flare series.*
6. POSSIBLE COMMON RELATIONSHIP OF TIlE ACTIVE REGIONS TO k IONG IVEI S'B-PIIOTOSPIIERIC SOURCE OF MAGNETIC FLUX
The history of the rise and decline of Plage 15031 and 13266 given in Sections3 and 4, forms part of a larger study made by the author of the waxing- and waningsof all active regions with a zone of northern solar hemisphere extending betweenheliographic Carrington longitudes L25
° and L85', over the periods of their vis-
ibility from Earth during Carrington Rotation 1659-1672 (i.e. within the bridginginterval of a terrestrial year from September 1977 to September 1978).
Against this background, which involved the study of 30 independent plages,some 8 of which executed three or more traversals of the solar disk, it was notedthat the ability of a particular plage to survive depended directly on the level ofcomplexity of Its underlying magnetic field. Within the atmospheric zone considered,strong fields surfaced only within a narrow band extending roughly from mean lat-itudes Ni6' - N26'. To the north and south of this domain, the fields that emergedwere markedly weaker with no recorded births to the north of mean latitude N3V andto the south of mean N130 .
Since the level of flaring observed within a particularplage was itself a close function of the level of complexity of the underlying mag-netic field, the locations at which energetically important flares were triggered
Compare with the recovery of the magnetic field strength exhibited on twooccasions by that part of Mt. Wilson Group 14284 directly underlying the triggerpoint of a sequence of three homologous proton flareq In 11c ath Plage 5265,(OcKenna-Lawlor, 1981).
133
showed the same latitude dependency as that exhibited by long-lived plages andcomplex magnetic fields.
Those plage sequences producing the various proton flares considered (and nan>vother energetically significant events as monitored by their electromagnetic accom-paniments), were further spatially associated with a persistent, although somewhatintermittently yielding "well-spring" of magnetic flux, situated between heliographiclongitudes L52
° - L79**. Specifically, Mt. Wilson Group 19984, underlying "November"
Plage 15031, was reported on emergence** to be located at mean latitude N23' andmean heliographic longitude L61
°.
Ht. Wilson Groups 20018 and 20019, underlying that western part of Plige 13266associated with the production of the 'pril-May proton flare, (see Section 4) were,on the other hand, reported on emergcnce to extend between mean latitudes N|9O 25'and mean heliographic longitudes L71
0-L79*. A temporally and spatially related
group which developed at mean latitude N25* and mean heliographic longitude L52was especially associated with energetic flaring showing a different morpholog tothat of the homologous proton events but present in the eastern part of the sameplage.
These spatial relationships suggest that while their component plage struc-tures and precise positions were different, the important proton flares of November22, 1977 and April-May 1978 may have constituted responses to instabilities as-sociated with the presence beneath the photosphere of an anomalous region cap:.hleof causing the episodic emergence over an extended time period of strong magneticflux within the confines of a solar zone lying roughly between latitudes Ni'
° - 26'
and mean heliographic lngitudes L52°
- L79°.
Compare this finding with the assertion by Svestka (1968) that 'proton fliareactivity regions are not randomly distributed on the solar disk but tend to occurin complexes of activity which stay on the solar surface for many months and evenyears". See also Dodson and }{edeman (1968) and a general discussion of related workby Svestka (1976).
Acknowledgments
:ly best thanks are due to E. Ruth Iledeman of the ,Ic~tath-Hulbert Obseratorv forvaluable discussions and for the gift of on and off-band Mcllath H-alpha spectro-heliograms of the proton flares of April 28 - May 1, 1978.
* Another "well-spring" of flux, associated with the production, in other plagesequences, of energetically significant flares between September !977 and September1978, were located, within the same latitude limits, between heli(;,raphic longtitudes129
° and 1,32'.
** A systematic change with time of solar longitudes in a direction opposite tothat of solar rotation was exhibited by all of the plages studied.
134
References
Dodson, HI.W., and lledeman, E.R.: 1968, "Some patterns In the developnent of centresof solar activity, 1962-19b8", in K.O. Kiepenheuer (ed.), Structure andDevelopment of Solar Active Regions, Proc. IAU Symp. No. 35, D. Reidel Publ.Co., Dordrecht, p. 5b.
IcKenna-Lawlor, S.I.P.: 1980, "Short-term prediction of the potential of an activere :ion to produce recurrent proton flares", in R.F. Donnelly (ed.), Solar-Trrrest rijd Prodi.'tion ProceedIings, Vol. III, U.S. Department of Commerce, NOAA,EI.M, Ktui.1er, Cotorado, p. C-12.
.vvetki, *'.: 1968, "Loop prominence systems and proton-flare active regions", inK.). Kiopenheoer (ed.), structure and Development of Sollar Active Regions,Proc. I \L Svnp. No. 35, D. Reidel Publ. Co. , Dordrecht, p. 287.
wv tk., '.: 1970, Solar Flares, Geophysics and Astrophysics 'ounograms Vol. 8, D."ei.,,I Pohl. 'o. , Dordrecht.
Discussion
'!)rvcr: o c rnin,, the two pairs of flares (28, 29, 30 April 978 and I lay 1978),,I; von think that .i repetitive twist-flare-relaxation, followed by a repeti-
ti,,n , this ;attern (such as the Ideas of Tanaka and Nakagawa) could have taken:'LI " hier'
'Icenna-Lawt, r: That is very interesting sogestinn. It is certainly feasible andw,+ tI k , t, , expl,re the po,;thilttv further.
135
0!
0
11. Theoretical MHD Simulations ofCoronal Transients and
Interplanetary Observations
M. lirner
Space ,nvironment Labo)rator. NOA,-A/ERL
loulder. Colorado 80303. ( .S.A.
R. S. SteimlollfsThe I ni ersit% o Xlalamnia in llunts~ille
Ihulsmille. .labania 35899. 1 .S. A.
Z. K. S1161h
Space lnvironnment Laborator%. NOAA/IERLBoulder, C)lorado 80303, U.S.A.
Abstract
A major lonq-ranqe goal of theoretical simulations of solar-aenerated distur-bances (transients, coronal holes, etc.) is the realistic modelinq of a propa-qating disturbance from the sun into and throuqhout interplanetary space. Simu-lations of this kind, using MHD fluid theory, must always be confronted withobservations in order to assess the deqree to which one or the other isinadequate. We describe some of this on-goinq work which is concerned with hothone- and two-dimensional, time-dependent MHn simulations.
The first example of such work begins with the simulation ofa flare-producedcoronal transient. In this case, a sudden current pulse is assumed to produceemerging magnetic flux. This "magnetic pulse" is sufficient to drive a weak shockwave into the solar wind by virtue of an outwardly-directed Lorentz force.Explicitly the pulse is assumed to consist of a !inearly-increasinq (with time)maqnetic field of 0.72 G at the base of a closed maonetic topolooy (initially inhydrostatic equilibrium) to a value ten times larqer over a In-minute period. Theshock achieves a velocity of 230 km 's- (which would he superimposed upon theexisting solar wind).
137
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k -- *E
A second example is concerned with a series of corotatinq interaction regionswhich were observed during a 60-day period by Pioneers 1n and 11 in ]071 prior to,during, and subsequent to the former spacecraft's flyby of Jupiter. Anopportunity for a stringent test of our one-dimensional model was made possible bythe nearly radial alignment of these two spacecraft. By using the upstream space-craft data as input into the time-dependent MHP model, we obtained a simulation ofthe solar wind characteristics at Jupiter that compared very well to the actualsclar wind plasma and magnetic field data that were measured by Pioneer in at thatplanet.
These examples dramatically illustrate the need for extensive observationalinformation at the inner boundary in order to assess the validity of any compre-hensive, self-consistent theoretical model. The second example dramaticallyillustrates the desirability of using actual data for both input and downstreamcomparison in order to assess the validity of a relatively straight-forward, yetcomprehensive, self-consistent theoretical model. On the other hand, the firstexample has not been tested with actual data as yet, although remote observationsof white-light density enhancements and type II shock velocities offer ouidancefor post-flare diagnostics. In the absence of additional data close to the sun(such as densities, velocities, temperatures and magnetic fields), we suanest thatthere is, at this time, no alternative to resorting to certain assumptions (suchas the current pulse used herein) concerninq the forcing function for initiationof multi-dimensional MHP simulation models of cause-and-effect.
I. INTRODUCTION
it has become abundantly clear to observers and theoreticians alike that thefield of solar-terrestrial physics will progress only through coordinated, inter-disciplinary observations suggested by the title of this symposium and theirdirect confrontation with theoretical models and simulations. The ambitious plansoutlined for the spacecraft partners in the observational phase in the lQRPs andIQQOs by Rohlin and Chipman (lQO) should be complemented by a strong theoreticalprogram.
It is our contention that dynamic MHD models and simulations should play amajor role in such an international program. Non-self-consistent phenomenologicalmodels (often referred to as "cartoons" in the magnetospheric physics community)are no longer adequate or satisfactory as noted by the U.S. Space Science Board(1978). Kuperus (1980) notes that numerical fluid calculations give promisingresults, notwithstanding the restrictions the theoretician is forced to make. Henotes, further, that these calculations ought to be made in three dimensions usinokinetic theory instead of the fluid approximation in order "to give full credit toall possible wave modes and interactions." Indeed, the philosophy and strategyfor the temporal three-dimensional approach, taking the solar surface and theboundary of the heliosphere (as the inner and outer boundaries, respectively) intoaccount, have been outlined by Wu (1q8n) and Yeh (1Q80). Also, Nakagawa, Wu andHan (1980) have described a quasi-three-dimensional MHP model (often referred toas the 2 1/2-dimensional model in fusion research). In addition, both Wu (1QRn)and Cuperman (1980) have noted the relative merits and limitations of the fluidvis-a-vis kinetic approach and the need for more physical insight via a "hybrid"
treatment. This latter approach has met with some success in pioneering simu-lations for laboratory fusion processes (Liewer and Krall, 1q73). Cuperman (IQ8O)
138
recommends a three-pronged strategy in this regard: (i) development and use ofhigher-order moments (fluid) equations as well as more realistic closureconditions and transport coefficients for the macroscopic description of the solarwind; (ii) use of computer simulation experiments on the nonlinear collectiverelaxation process through particle-wave-particle interaction due to plasmaelectromagnetic instabilities; and (iii) the hybrid strategy wherein collectiveinteractions in the evaluation of transport coefficients, as deduced from quasi-linear theory and computer simulation experiments, are incorporated into thehigher order moment equations.
The reader will recognize that an ambitious plan, such as that outlinedabove, should advance in an orderly fashion from the present state-of-art. Thepresent models contain the basic physics necessary to study the large-scale, grossobservations but do not contain the mathematical and physical elegance notedabove. Such elegance, we believe, is required for study of the finer details ofthe bulk flows and are not necessary for our present purposes in the discussionwhich follows below. This paper will direct attention to several examples repre-sentative of present dynamical, MHD, fluid models. We have chosen the followingexamples:
(i) Two-dimensional, time-dependent, planar MHO model for coronal massmotion as a result of an emerging magnetic flux through the ohotosphere; and
(ii) One-dimensional, time-dependent, MHD model for solar wind dynamics asa result of interaction among multiple, coronal-hole-generated, corotatingstructures which were measured by Pioneers 10 and 11 during two solarrotations when these spacecraft were essentially co-linear with the sun.
The strategy adopted in the first example has been reviewed most recently byWu (1Q8O) and extended by Steinolfson and Wu (1QRn.) It has heen recognized bysome observers (c.f., Anzer, IQO; Sheridan, 1QRO, Gergely and Kundu, 1QPO; andMaxwell and fryer, 10O) as a viable model for providing insight to the phenomenonof coronal transients (c.f., Michels et al., IQ8O). It is important to note thatthis approach does :. consider the original flash phase of a flare or triggeringmechanism for an eruptive prominence. This primary phase is associated withcurrent sheet disruptions (Syrovatskii and Somov, 1QP08 Somov and Syrovatskii,1Q80; and Sermulina et al., iq8O) and rapid conversion of the magnetic energystored within force-free magnetic topologies into thermal and kinetic energeticforms. The subsequent phase does not consider these initial physical processes;instead, their consequences may be considered as producing a thermal or magneticpulse with a certain magnitude as well as temporal and spatial extent. uringthis secondary phase (coronal mass motion which moves, generally preceded hy an*MHD shock if the pulse is sufficiently large, into interplanetary space), therelative degree of wave propagation and mass motion is determined hy dynamicalinterplay between plasma motion, plasma forces, and magnetic forces ascharacterized by the plasma (ratio of thermal to magnetic pressure). "Magneticcontrol", then, is actually one of modulation of this secondary phase. That is,as 0 - 0, there will be minor plasma mass motion, hut there will he rapid MHnwave motion which will steepen into a weak fast-mode shock. Alternativelyas p increases to, say, 0.1 or greater, substantial mass motion (supplied from thechromospheric layer), accompanied by nonlinear MHn waves' steepening into a shockwave, most definitely will take place and will he detected as a type 11 radioburst. Additional diagnostics will include type IV radio emission due,presumably, to qyrosynchrotron radiation (or even plasma oscillations as sunaestedfor very bright events by Duncan, Stewart and Nelson, 1080) and white lightcoronal transients. The case discussed here is the simulation of the coronalresponse to a magnetic pulse. It will be shown that p (r, A. t: where r is radialdistance, A is the solar latitude, and t is time after the pulse is initiated)
139
exhibits values over a wide range, both > 1. Thus, the degree of magnetic modu-lation ("magnetic control," if you wi<l) must be examined to discern itsdependence on the relative contribution of inertial, thermal, gravitational, andmagnetic forces.
The strategy adopted in the second example noted above is to consider theextension to 5 AU of a series of high speed streams during a 60-day period in late1973. Coronal-hole-generated high speed streams, rather than flares, are assumedto he responsible for these observations by Pioneers In and 11. The basic physicsof the "secondary phase" must be the same as for that following flares, eventhough the coronal hole's solar surface boundary conditions and their initialvalues are different - and undoubtedly milder - than those for the flare so thatthe ensuing mass motion and wave propagation are also milder. However, it iswell-known that the longer-lasting energy and momentum addition to the solar windwill also produce MHD shock waves at large distances (sometimes even within 1AU). It is now known that these convective pressure pulses, sometimes bounded byforward and reverse MHD shock waves, buffet the planetary magnetospheres of Earthand Jupiter and produce substantial, albeit temporary, deformation of their bowshock waves and maqnetopauses. The shocks themselves also produce local acceler-ation of solar wind protons to near-MeV energies as well as a chopper-like modu-lation of energetic Jovian electrons.
It is our intention, then, to present two examples of dynamic MHP modelinn -one close to the sun, the second far from the sun - to demonstrate the necessityfor modeling which considers the region of space from the sun outward, in a con-tinuous fashion, to the distant reaches of the heliosphere.
2. RESULTS AND DISCUSSION
We now discuss the two examples which were outlined in the Introduction.
2.1 Two-l)ilneisional, Time-I)epenrdent, Planar MIII) MNIoeI Near the Sti
The basic equations and assumptions made in the development of this modelhave been described by Steinolfson et al., (1Q7P) and will not he repeated here.Although thermal pressure pulses have mainly been considered (Steinolfson et al.,1978; Dryer et al., 1979; Dryer and Maxwell, 1979) the present example assumesthat the pulse at the coronal base is produced by emerging magnetic flux due to ahypothesized sub-photospheric current pulse. This work, described by Steinolfsonet al. (1979), assumes the pre-pulse, steady-state, coronal atmosphere to he per-meated by a force free closed magnetic topology. We repeat some of the featuresdescribed by this latter paper.
We consider a model in which the driving force (the solar event) is createdwhen new magnetic flux due to photospheric currents emerges through the solarsurface and interacts with the pre-existing field. There is a orowinn body ofobservational data which supports this possibility for some forms of solaractivity. This emerqinq flux model has been speculated on for some time from a
140
theoretical viewpoint as well (see, e.g., Tur and Priest, 1078, and Syrovatskiiand Somov, 1980). The driving force created by the emerginq flux consists of anincrease in the magnetic pressure producing an outwardly directed Lorentz force.The thermal pressure also increases in the simulated solar event, but it increasesas a result of the emerging flux rather than independently as in previousstudies.
The model we use is applicable to a meridional plane and is similar to thatused by Steinolfson et al. (1978). We assume that the solar atmosphere can betreated as a single fluid with negligible dissipative effects. The necessarytime-dependent equations are given there and will not he repeated here.
The atmosphere is assumed initially to be In hydrostatic equilibrium and toobey a polytropic law with a polytropic index, y, of 1.1. The initial thermo-dynamic conditions at the base of thg atp1sphere (1.0 rp) are temperature T = 2.Ax 10 K and number density n = 3 x 10 cm- . The initial maonetic field is assumedto be force free and is calculated as discussed by Stelnolfson et al. (1Q7p, PaperI). A schematic of the initial field is shown in Figure 1. The magnetic fieldconfiguration at time t = 0 in this Fioure is referred to as a "closed-field"configuration in Paper I since simulated solar events at the solar surface nearthe equator produce transients which must propaqate across the field lines inorder to reach the outer corona. The magnitude of the initial magnetic field B istaken to be such that B 0.72 G at the solar base at the equator.
MAGNETIC FIELD LINIS
0 Miute, 16 Inutes 21 miutes
] II 1
-. i tl p.l tl i l+t '+L t IK . 1 I
Figure 1. The magnetic field lines initially and after 16 and 2P minutes.The equator is at the left-hand edge on each plot and the solar surface atthe base. The magnetic field is in the clockwise direction. Several repre-sentatives field lines, labeled 1 and 2 at t = 0 min., are also shown at t =16 min.
The emerging magnetic flux is assumed to be due to a photospheric linecurrent flowing into the plane of the paper in Figure I at the equator and 7000 km
141
below the solar surface. This current produces an emerging field in the samedirection as the initial field so the Lorentz force created when they interact isdirected outward. The photospheric current is increased linearly over a 10-minuteperiod, and its final maqnitude enhances the magnetic field at the solar surfaceat the equator by a factor of 10 over its initial value.
The time-dependent coronal response to the emerging flux is calculated bynumerically solving the time-dependent equations in Paper I. The solution issymmetric about the equator. Thus, the computations are performed, and theresults are presented, only for the region shown in Fiqure 1.
The magnetic field lines, initially and after 16 and 2P minutes of physicaltime, are shown in Figure 1. The outward motion of the field lines is indicatedby the displacement of the field lines labeled 1 and 2. Since the electricalconductivity is infinite in this study, fluid particles remain attached to thefield lines. As a result, the field line initially at the solar surface repre-sents the contact surface at later times and separates initial coronal materialfrom new material flowing upward through the coronal base. The contact surface isindicated by the field line with the short hatched lines across it in Figure 1.
A shock is formed ahead of the coronal disturhance as a natural consequenceof the non-linear MHD. Once the shock is formed it moves outward without reaardto the emerging flux. The shock has propagated beyond the region shown in Fioure1 for t = 2R minutes. The initial coronal material between the shock and thecontact surface has been heated, compressed and accelerated by the shock, whilethe new material behind the contact surface has been heated and compressed by theemerging flux. The Lorentz force is directed outward between the shock and thecontact surface and is directed inward below the contact surface. The plasmabeta (p) is greater than unity above the contact surface and less than unity belowit indicating the relative importance of the magnetic and thermal forces in thetwo regions.
I VI PA T '
T.T
I 1.4I
Figure 2. The temperature after 16 and 28 minutes referenced to the initialtemperature at each point. The equator is at the right-hand edoe and thesolar surface at the top.
The major difference between the present results for emerging magnetic fluxand the results obtained for a thermal pressure perturbation (as in Steinolfson et
142
rW
al., 1978; and Dryer et al., 197q) is the presence of the inward directed Lorentzforce below the contact surface. The Lorentz force tends to retard the outwardmotion of the new material, but the resultinq increase in magnetic and thermalpressure drives the material outward. The tendency for the Lorentz force to trapthe material behind the contact surface causes the emeraina field to heat andcompress it as shown on the next two figures which show carpet plots of the temn-erature and density, respectively, after 16 and 2P minutes. The values plotted inthe figures are the physical quantities normalized by the initial values at eachpoint. As shown in the figures, the temperature increases by a factor of about 2at most, and the density by a factor of about 4. The contact surface is Justahead of the temperature maximum in Figure 2 and just ahead of the density minimumin Figure 3. This hot, dense region near the solar surface may be responsible forproducing the X-rays often seen in solar events. For all simulations using athermal pressure perturbation, the temperature and density below the contact sur-face are anti-correlated thereby reducinq the possibility of producinq X-rays.
0. "P7,P T
Finure 3. The density after 16 and 2P minutes referenced to the initialdensity at each point. Orientation is the same as for Fioure 2.
There is only a slight increase in the thermodynamic quantities across theshock wave since the shock is weak and y is small. There is a relatively laroeincrease in the plasma velocity across the shock, however, as indicated in Fiaure4 which shows the plasma velocity vectors.
It is apparent that emernino maonetic flux will produce an interplanetarydisturbance. This is more obvious in Figure which shows the shock and contactsurface trajectories alonq the equator. Althouah we have only followed the dis-turbance for 30 minutes, both the shock and the contact surface have reached con-stant velocities and will presumably continue on outward into the interplanetarymedium. For the thermal pressure perturbation used in the closed-field confinur-ation by Steinolfson et al. (107P), none of the material below the contact
143
PLASMA VELOCITY VICTORS
16 Minutes 28 Minutes
1.4 1.4 t I f r I I I I I I
a SI ,1 1, /, , ,,,
" - II I111//t//,
* i,,,, Shock* 171///,
C It a/ tl -.tt --.Suntacte I / , A /u • / , - - - - III / '
0 1 0Latitude, Oeg. Latitude, Deg.
Figure 4. The plasma velocity vectors after 16 and 2P minutes. Orientationis the same as for Figure 1. Each vector points in the direction of thevelocity and has a length proportional to the velocity at its hase. As n tedin the box for t = 16 mi., the larest velocity at this time is 12. km s-
surface reached the outer corona. Note that the shock velocity is uite low, muchless than the values of 1000 km/sec, or even qreater, detected by type 11 radioemission (Stewart, 1980). We have not investigated this example further to deter-mine the behavior of such a simple emerqinq flux assumption on a lonoer time scaleor by using large current pulse magnitudes. It might even he anpropriite tosuggest that flares with flash phases (e.q., Somov and Syrovat:' IQ Q0,Sermulina et al., 1980) associated with current sheet disruptions (not consideredhere) may well he shown, with such a model, to have large thermal pulses as well,as suggested by Dryer and Maxwell (1)79) and Maxwell and Dryer (IQO). It is ofparticular importance that cause-and-effect (the former as just discussed; thelatter as measured in the interplanetary medium) he studied with synoptic space-craft measurements. A limited amount of research in this reoard has beenperformed by Gosling et al. (1975) and Wu et al. (1976) who make phenomenoloaicalassociations of coronal transients, as observed via Skylab and ns0-7 white-lichtobservations, with Pioneer 9 and 10 (respectively) in situ plasma observations ofthe shock waves and disturbed solar wind which followed these shock precursors.
144
• 230 km s
Shock
% 110 km
Cont:,ct Surface
10 IS 20 25
Time, Minutes
Figure 5. Trajectories of the shock and the contact surface alona theeauator.
2.2 One-)irnensional, Thne-Dependent I IIID lodeil Far From the Sun: Direct Comparison with Pioneer]O/I1 Observations
The radial alignment of Pioneers 11 and 10 (at - 2.P and 5.n All,respectively) with the sun provided an excellent opportunity to test a one-dimen-sional MHD model (Steinolfson, Dryer and Nakagawa, q75) via a direct confronta-tinn with plasma and maqnetic field observations. The present study is an exten-sion of that already published by Dryer et al. (197R). It is an extension in timeto demonstrate the simulated corotatinn interaction reaions (CIRs) which buffetedJupiter before, during, and after the Pioneer 10 encounter with the Jovian magnet-osphere as discussed by Smith, Fillius, and Wolfe (1q78).
145
-Pioneer I1 Observations at 2.78 AU(2.28 Hour Averoges, Used for MHD Simulation)
700
v 600km/sec 500
65 rT
l i l i-
n 4
(cm-3) 32 ,
0
3(0 320 330 340 350 3601973 (Days)
Figure 6. Pioneer 11 ohservations at 2.7P All. These parameters are used asinput for the 6H! modelinq.
Fiqure 6 shows the solar wind velocity, density, proton temperature, andazimuthal magnetic field measured from flay 3nl1 to flay 361, 1q73, hy Pioneer 1!.This spacecraft was initially located at -. 2.7P All and was aligned with Pioneer ilOwhich (during this epoch) approached Jupiter, flew past this planet within itsmannetosphere, and continued heyond Jupiter in the solar wind. The data shown inFinure 6 were used as a 60O-day "forcina function" for the one-dimensional model
146
"-~~~~~1 .. .. .. .. I I II l . .l i . . . . .
Pioneer 10 Observationsat 5.02 AU (2.28 Hour Averages)
----MHD 1-D Simulation
V Pioneer 10(k/sc Observation Jupiter453sec Encouniter 550
350 -... 450
Simulatin
n 12, Observation
_c3) 08!04~
Simulation - 108
04
T 10' 10(o K) 103 Observation10
- Observation
-21 Simulation -~ 1
310 320 330 3-4-0350 3----- 2'0Days (1973)
Figure 7. Pioneer 10 observations and the MRP model simulations at S All.
147
noted above. The time-dependent simulation provided temporal and spatial computa-tions for the ensuinq response which was tracked beyond 5 AU. The time scale ofthe simulation obtained at Jupiter, at 5.02 AU, was corrected for the variationscaused by the changing spacecraft separation during the 60-day observationperiod. Also, the spacecraft velocities were taken into account. This proceduremade it possible to "pull out" the simulated parameters as "observed" by hypo-thetical instruments in front of the planet's bow shock wave. It was, therefore,possible to make a direct comparison with Pioneer 10's actual observations beforeand after its penetration of the Jovian magnetosphere. It was also possible topredict the varying solar wind's pressure impingement upon the Jovian maoneto-sphere during the time of the encounter.
Pioneer 10 Observations at 5.02 AU(2.27 Hour Averages)Pioneer II Simulation
S --
I I I . .. . .. . T I I T T I I 1 T I 1z
E I PioneerlO I10' Encounter
1610°F10 Observation
S1610-8
0 10-1
Simulation - -0-1
310 320 330 340 350 3601973 (Days)
Figure R. Corotating interaction regions at 5 AU as defined by theconvective pressure.
This comparison is shown in Figure 7. It is clear that comparison of thephasing and magnitudes of the individual parameters is satisfactory. The CIRs,themselves, are nore clearly discerned in Figure 8 which shows the convectivepressure, dyn cm'1. The simulated and observed pressure pulses agree to within 3hours with two exceptions: the first pulse and the partial pulse which were ob-served as the spacecraft emerged from the Jovian magnetosphere. We conclude thatthe MHD fluid approximation, even confined to one-dimension (the radial one) andone-fluid, is highly satisfactory for demonstrating the essential physics ofdynamic solar wind streams.
148
Acknowledgments
The work performed by RSS was supported through NOAA Grant Mo. NARORAPODOI2.We wish to thank Dr. T. Yeh for his suggestions during the Dreparation of thisreport.
References
Anzer, I.: 19PO, "MHD aspects of coronal transients," in M. Dryer and E.Tandberg-Hanssen (eds.), SoZar and Interplanetary Dynamics, Proc. IAI) SymposiumNo. Q1, P. Reidel Publ. Co., Dordrecht-Holland, p. 263.
Bohlin, J. D. and Chipman, E. G.: 1980, "A program for the observations of the sunand heliosphere from space 19R0-IOQ5," in M. Dryer and E. Tandberq-Hanssen(eds.), solar and Interplanetary Dynamics, Proc. IAU Symposium No. 91, D. ReieelPubl. Co., Dordrecht-Holland, p. 523.
Cuperman, S.: 1980, "Physical processes and models of interplanetary responses:suggested theoretical studies," in M. Dryer and E. Tandbero-Hanssen (eds.),Scinar and interpZanetary Dynamics, Proc. IAU Symposium No. 0I, ). Reidel Puhl.Co., Pordrecht-Holland, p. 459.
Dryer, M. and Maxwell, A.: 1979, "Radio data and a theoretical model for thefast-mode MHD shock wave generated by the solar flare of 1Q73 September q, 18:26LIT," Ast~orhhe. J., 231, 945-959.
Dryer, M., Smith, Z. K., Smith, E. J., Mihalov, J. p., Wolfe, J. H., Steinolfson,P. S., and Wu, S. T.: 107R, "Dynamic MHn modelina of solar wind corotatinnstream interaction regions observed by Pioneer 10 and 11," J. Geophys. Res., 83,4347-4352.
Dryer, M., Wu, S. T., Steinolfson, R. S. and Wilson, P. M.: 1)7q, "Maqnetohvdro-dynamic models of coronal transients in the meridional plane. II. Simulationof the coronal transient of 1973 August 21," Astrophzys. J., 227, In59-1071.
Duncan, P. A., Stewart, P. T., and Nelson, G. J.: 1980, "Recent very hricht typeIV solar metrewave radio emissions," in M. Dryer and E. Tandberq-Hanssen (eds.),oZir and InterpLanetary Vynamice, Proc. IAU Symposium No. 11, D. Reidel Publ.
Co., Dordrecht-Holland, p. 381.
Gerqely, T. E. and Kundu, M. R.: 1980, "Decameter radio and white light obser-vations of the 21 August 1973 coronal transient," in M. Dryer and F. Tandhera-Hanssen (eds.), Solar and Interplanetary Dynamics, Proc. IAU Symposium No. 91,P. Reidel Publ. Co., Dordrecht-Holland, p. 245.
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References
Goslinq, J. T., Hildner, E., MacOueen, R. M., Munro, R. H., Poland, A. I., andRoss, C. L.: 1975, "Direct observations of a flare related coronal and solarwind disturbance," Solar Phys., 40, 439-448.
Kuperus, M.: 1980, "Symposium summary," in M. Dryer and E. Tandbero-Hanssen(eds.), Solar and Interplanetary Dynamics, Proc. IAU Symposium No. 91, h. ReidelPubl. Co., Dordrecht-Holland, p. 547.
Liewer, P. C. and Krall, N. A.: 1973, "Self-consistent approach to anomalousresistivity applied to theta pinch experiments," Phys. Fluids, 16, 1053-1963.
Maxwell, A. and Dryer, M.: 1980, "Radio data and computer simulations qeneratedby solar flares," in M. Dryer and F. Tandberq-Hanssen (eds.), Solar and Inter-planetary Dynamics, Proc. IAU Symposium No. 91, P. Reidel Publ. Co., Dordrecht-Holland, p. 251.
Michels, D. J., Howard, R. A., Koomen, M. J., Sheeley, N. R., Jr., andRompolt, B.: 1QPO, "The solar mass ejection of 8 May 1979," In M. Dryer and E.Tandberg-Hanssen (eds.), Solar and Interplanetary Dynamics, Proc. IAU SymposiumNo. 91, D. Reidel Publ. Co., Dordrecht-Holland, p. 387.
Nakaqawa, Y., Wu, S. T. and Han, S. M.: 1980, "Dynamics of coronal transients:two-dimensional non-plane MHD models," in M. Dryer and E. Tandberg-Hanssen(eds.), S, ' ad Interplanetary Dynamics, Proc. IAU1 Symposium No. 91, P. ReidelPubl. Co., Dordrecht-Holland, p. 495.
Sermulina, B. J., Somov, S. V., Spektor, A. R. and Syrovatskii, S. I.:1980, "Gasdynamics of impulsie heated solar plasma," in M. Dryer and F. Tand-berg-Hanssen (eds.), So~ar and Interplanetary DYnamics, Proc. IAU Symposium No.91, D. Reidel Publ. Co., Cordrecht-Holland, p. 491.
Sheridan, K. V.: 1980, "Evidence for open field lines from active regions:short communication," in M. Dryer and E. Tandberq-Hanssen (eds.), Solar andInterplanetary. Dynamics, Proc. IAU Symposium No. 91, D. Reidel Publ. Co., Dord-recht-Holland, p. 261.
Smith, E. J., Fillius, R. W. and Wolfe, J. H.: 1978, "Compression ofJupiter's magnetosphere by the solar wind," J. Geophys. Res., 63, 4733-4742.
Somov, B. V. and Syrovatskil, S. 1.: 1980, "Magnetically driven motions in solarcorona," in M. Dryer and E. Tandberg-Hanssen (eds.), Solar and Interplanetary.Dynamics, Proc. IA11 Symposium No. 91, D. Reidel Publ. Co., nordrecrt-Holland,p. 487.
Space Science Board: 1978, "Space Plasma Physics: The Study of Solar SystemPlasmas, 1978, - Colgate Report," Vol. I., National Academy of Sciences/NationalResearch Council, Washington, D.C., pp. 12-13.
Steinolfson, R. S., Dryer, M., and NakaGawa, Y.: 1975, "Numerical MHDsimulation of interplanetary shock pairs," J. Geophys. Res., 80, 1223-1231.
Steinolfson, R. S. and Wu, S. T.: 1980, "Evolution of coronal maqneticstructures," in M. Dryer and E. Tandberq-Hanssen (eds.), Solar andInterplanetary Dynamics, Proc. IAU Symposium No. 91, D. Reldel Puhl. Co., Dord-recht-Holland, p. 483.
150
References
Steinolfson, R. S., Wu. S. T., Dryer, M. and Tandberg-Hanssen, E.: 197R,"Magnetohydrodynamic models of coronal transients in the meridional plane. I.The effect of the magnetic field," Astv'ophys. J., 219, 324-335.
Steinolfson, R. S., Wu, S. T., Dryer, M. and Tandberq-Hanssen, E.: 1Q7Q, "Effectsof emerging magnetic flux on the solar corona," in H. Rosenbauer (ed.), Proc.Solar Wind IV Conference, Springer-Verlag, Heidelbura, in press.
Stewart, R. T.: 1980, "Transient disturbances of the outer corona," in m. nryerand E. Tandberg-Hanssen (eds.), Solar and Interplanetary Dynamics, Proc. IAUSymposium No. 91, D. Reidel Publ. Co., Dordrecht-Holland, P. 333.
Syrovatskii, S. I. and Somov, B. V.: 19RO, "Physical drivina forces and modelsof coronal responses," in M. Dryer and E. Tandberg-Hanssen (eds.), Solar andInterplanetary Dynamics, Proc. IAII Symposium No. 91, D. Reidel Publ. Co., ford-recht-Holland, p. 425.
Tur, T. J., and Priest, E. R.,: 1978, "A trigger mechanism for the emeroina fluxmodel of solar flares," Solar Phys., 58, 181-200.
Wu, S. T.: 1980, "Theoretical interpretation of travelling interplanetary phen-omena and their solar origins," in M. Dryer and E. Tandberq-Hanssen (eds.),Solar and Interplanetary Dynamics, Proc. IAU Symposium No. 91, D. Reidel Publ.Co., Dordrecht-Holland, p. 443.
Wu, S. T., Dryer, M., and Han, S. M.: 1976, "Interplanetary disturbances in thesolar wind produced by density, temperature, or velocity pulses at 0.Op AU,"Solar Phys., 49, IR7-204.
Yeh, T.: 1980, "Global modeling of disturbances in the corona-interplanetaryspace," in M. Dryer and E. Tandberg-Hanssen (eds.), Solar and InterplanetaryDynamics, Proc. IAU Symposium No. Q, D. Reidel Publ. Co., Dordrecht-Holland,p. 503.
151
12. Three-Dimensional Distributionof the Solar Wind Velocity
Deduced From IPS Observations
T. Kakinuma, H. Washimi, and M. KojimaThe Research Institute of Atmospherics
Nagoya UniversityNagoya, Japan
Abstract
-The tPS observations of the high-latitude radio source 3C48 are used to erivethe distribution of the solar wind velocity in lonitude and latitude for periodsof several rotations during 1974-1979, by model calculations. The observationsduring 1974-1978 show the existence of the north polar high-speed (800 km/sec)rcgion. This region extended from the pole down tq 45 - 50 latitude for mostof 1974-1977, while the region contracted up to 80 in 1978 and was hardly observedin early 1979.
I. INTRODUCTION
One of the important subjects in the ohservation of the Jnterplanetary scin-tillation ([PS) of radio sources is to derive a three-dimensional picture ot the so-
lar wind. Dennison and Hewish (1967) first made three-station observations of IPS
153
and found that the solar wind at high latitudes was faster than at low latitudes.Watinabe et al. (1974) investigated the relation between the solar wind velocityand EUV brightness of the corona. They found an inverse correlation between them
at high latitudes. This explains the negative results obtained for the latitudedependence of wind velocity, observed by hlewish and Symonds (1969) and also byVitkevitch and Vlasov (1972), and suggests the change of wind velocity at high
latitudes is due to solar activity. Coles and Rickett (1976) showed that theaverage solar wind speed increased away from the equatorial plane with an averageaverage gradient of 2.1 krm/sec per degree of latitude during 1971-1975, and recentlythey reported a solar cycle change of the high-latitude solar wind (Rickett andColes, 1979). Kakinuma (1977) derived the latitudinal distribution of the solarwind velocity in 1974 by a model calculation and showed that the solar wind velocitywas 800 km/sec at latitudes above 45'.
In order to derive the velocity distribution in longitude and latitude, it is
necessary to observe, simultaneously, many radio sources in different directions.The number of radio sources which are now observed is not sufficient to cover allsolar coordinates at one time. However, when the corona is stable during severalsolar rotations, we can infer the large-scale three-dimensional structure of the
solar wind from observations of a small number of radio sources, even from observa-tions of one high-latitude source. Sime and Rickett (1978) derived the velocitydistribution in longitude and latitude for each year from 1973 to 1977 by assuming
that the solar wind speed at the point of closest approach of the line of sight tothe sun was measured by the IPS method and by averaging together IPS observationsof several radio sources over several rotations. Kakinuma et al. (1979) alsoderived the latitude and longitude distribution of the solar wind speed in 1974 by
model calculations using the observations of 3C48. In this paper, we discuss the
velocity distribution from 1974 through 1979.
2. MODEL CALCULATION
i'S observations at Nagoya University are being carried out at three stations:Toyokawa, Fuji, and Sugadaira. The observing frequency is 69 TtHz. Details of our
receiving system have been described by Kojima et al. (1979). The velocity of thediffraction pattern drifting across the earth is derived by cross-correlation
analysis (Kakinuma et al., 1973). We calculate three cross-correlation functionsof intensities between pairs of stations and then calculate the velocity from the
time lags of the cross-correlation functions under the assumption of a radial flowand a frozen pattern. We take into account the anisotropy of the pattern in the
c alculation.
Tiansen et al. (1976) have shown that the coronal structure remains stableover periods of several months during the declining and the minimum phases ofsolar activity. In such a case, we can infer the distribution of the solar windvelocity in longitude and latitude from IPS observations of a high-latitude source.The IPS pattern speed is a weighted average of the transverse components of thesolar wind velocity along the line of sight to the radio source. The heliographiccoordinates of sources of the solar wind crossing the line of sight change everyday with the movement of the radio source relative to the sun and to the solarrotation. Thus observations of the pattern speed during several rotations may beconsidered to include information about the velocity distribution over a widerange of longitude and latitude. We have attempted to derive the velocity distri-bution of the solar wind by model calculations using the observations of 3C49 fromwhich we can obtain a high signal-to-noise ratio.
154
Our model consists of two distributions: one for the "background solar wind"which corresponds to the lower envelope of the daily plots of tPS pattern speed,and the other for localized high-speed streams which correspond to enhancementsabove the lower envelope. The observations of 3C48 in 1974, for example, showthat this model is reasonable (Kakinuma, 1977). In the calculation, we used theweak scattering and weak scintillation approximation. We assumed a power-lawspectrum (power-law index = -3) of electron density fluctuation ( A ne) theinverse square dependence on the radial distance of electron density fluctuation,and a Gaussian brightness distribution of the radio source. We used the valuesgiven by Readhead and Hewish (1974) for source size. The proton flux is commonlylewish (1974) for Sourcesize. The proton flux is commonly assumed to be nearlyconstant; thus An e . if Ane - n . However, if we assume Ane V
1
a velocity more than I00 km/se i 11 1 e required at the pole to predict the lowerenvelope of the observations in 1974. A comparison of spacecraft and IPS measure-ments of the solar wind (Coles et al5, 197Pl) also shows that Ane is not propor-tional to n . Thus we assumed (Ane )- V , except for the region of stream-stream inte action at the leading edge of the stream. The stream model used inthe calculation is shown in Figure 1. In this model, An. at the leading edge, is
times higher than the ambient level at a distance of I AU from the sun in theequatorial plane, and decreases with increasing latitude and with decreasing radialdistance. We used Carrington's synodic rotation rate for all latitudes, and alsoassumed that there was no radial distance dependence of the solar wind velocity.
Using the sane method as used in the daily analysis of observations, crosspowerspectrum from 0.1 Iliz to I liz, and then the intensity cross-correlation function bythe Fourier transformation, the pattern speed was calculated from the time lag ofthe cross-correlation function. First we determined the distribution fit to thelower envelope of daily plots of pattern speed and then included the enhancementsdue to localized high-speed streams.
3. VELOCITY DISTRII1TION
A comparison of the observed and the predicted IPS pattern speeds for February-July 1974 is shown in Figure 2. There are some differences between them, as we haveneglected the temporal variations of velocity distribution, but this figure showsthat our model is a good approximation. In our model, in the northern Hemisphere(above -7' latitude), the speed of the "background solar wind" increases with lat-itude, and can be given by 300 + 500 x sin2 (2 x latitude) km/sec up to 45 , and 800km/sec at latitudes above 45 as shown in Figure 1. During the time period Feb-ruary-July 1974, there were two corotating high-speed streams. One was centeredon 250 Carrington longitude and was an equat~rward extension of the north polarhigh-speed region. The other, centered on 90 Carrington longitude, was an exten-sion over the equator of the southern high-speed streams, which was most likely an
equatorward extension of the south polar high speed region. The existence of thehigh-speed stream associated with an equatorward extension of the solar coronalhole has been pointed out by many authors. This distributon is consistent with theobservation of the K-corona (ansen et al., 1976) and the averaged distributionderived by Sime and Rickett (1978). This distribution also shows that there is nolocalized high-speed stream at latitudes above 45', agreeing with louminer's re-sult. lie has found by scintillation index measurements that there is no stream-stream interaction at latitudes above 40 (Iouniner, 1977). The enhancementq inpattern speed in high-latitude observations are due to high-speed streams crossingthe line of sight at latitudes below 45 (Kakinuma, 1977); speed enhancements inhigh-latitude observations do not always mean the existance of high-speed streamsat high latitudes.
155
CARRINGTON LON6ITUDE0* W.90 180" 270" 360.
A0.0 -A----0.1
0,5 -x
0.S STREAM N STREAM
1.1.0 i I7- 0. wC" 180" 270" 360"
V
2.0-
10 - N__ _ _ _ 300_ __ __
000 - 0 tI i
01 30' 60" 90"LON6ITUDE HELIOGRAPIC LATITUDE
Figure 1. The stream model and the model of latitudinaldistribution of the solar wind velocity for 1974.
1974 (3C 49)
.. 9DEL (ALCULA'104
OBSERVATION
KWSEC
40 so 60 70 SO 90 1.0 l 120 110 10 I0 60 170 10DAY Of YAR 1974
Figure 2. Observed and calculated IPS pattern speeds andthe model calculations (using an 800 km/sec source region)for 1974. The IPS source is 3C48.
156
19A X0 48)
1WO
I II I I L ALUA I I
4o0
Q0 so 0 1 0 so 90 100 110 120 130 140 1~ 1)0 177 I .DAY OF AAR 1974.
Fi:iiro 3. Observed and calcuilated IPS pattern speeds andthe model calculatioas Cisin an 800 km/sec tilted polar
s e h',e d s o ir c r~ e'o n) f or 1974. The II'S sourve is 3C 4 . n.
1974 0CI11
%730L ALJLAIO
110 12 .IS•. 1 IO ID I w "......
f I OS'
103
50 .70 g go 1 i010 I I iiA K. ' I
1AY 4f 1tA, g
Frigure 4. Observed and calctilated IPS pattern qpeed- andth
th odel calcl~ations sn an 800 kin/se tilted polar p-pe
il. -edsource re;ion) for 1974. The U' S source, is 3 LCo.rain
10 120 3l 10 10 30 10 180 10 73 '0 .17 .0 .0
Figure 4. Oh.urved and calculated IIPS pattern speeds and the'model ralclatlons (using an 801) kin/sec ti It(,d polar high-speedsource regqion) for 1974. The I'Sq source is 3ClhlI (UC1)S observations3).
157
1975 (K298)
. . LeI CALCULATIN
V
1000
... .. . .. .. ... . ... ..200
0 I I I I I I I I I
220 230 20 0 260 270 280 290 300 310 320 530 09 50 %ODAY X YEAR 197
Figure 5. Observed and calculated IPS pattern speeds andthe model calculations (using an 800 km/sec source region)for 1975. The IPS source is 3C298.
1976 (3C 98)
MODEL CALCULATION
V 5*
600
200
40 SO 60 70 so 90 I90 110 120 130 140 IO 160 170 Ig)DAY Of YEAR 197t
Figure 6. Observed and calculated IPS pattern speeds and themodel calculations (using an 800 km/sec source region) for 1976.The [PS source is 3C48.
158
R; O 4S)
. . O1D.L ALLU.ATI
9BSL 9VAT IA
10 0 -0 5 0 ~ ~ 10 10 10 1F .C 1f 1, 7 5
&V 14, , .,' ": ... , ,,, ..
wo # - OF Co0I!
500 O .. . o* ,
40 w 1 1 10 I 1. E 1. E I ,l It t 111 !8
50 .0 5) 70 8) 0 IF 10 1 0 12F 10 1 50 I SO15 IV 1') 550
Figilre 7. Observed and calculated IPS pattern speeds and themodel calculations (usinp an 800 km/sec source region) for 1977.The [PS source is 3C48.
.159 AY NIAR4
Th0 IP soi c i0 3-8 . 0 l 2 lO 1 r
159
Sime and Rickett (1978) have conc Luded, from the averaged ve loci ty distr ib-tion, that the polar high-speed regions were not symmetric about the spin axi;, hua
were centered more than 300 from the rotation axis during 1973-1975. As shown in
Figure 3, the tilted polar high-speed regions seem to be a good approximation, but
the distribution shown in Figure 2 is much better for the northern hemisphere in
in 1974. UCSD observations for 30161 (southern high latitude source) in.1974, how-
ever, are well approximated by the polar high-speed regions tilted by 45 , as shown
in Figure 4. K-coronameter observations (Hansen et al., 1976) also show that thenorth polar hole extended from the pole to 45 latitude and was symmetric about therotation axis in 1974, while the south polar hole was displaced from the pole during
fron the pole during April-September 1974.
The calculations for 1976 (Figure 6) show that there were two corotating high-speed streams in the northern hemisphere which were equatorward extensions of the
polar high-speed region: one centered on 240 Carrington longitude, was a stablewide stream and the other, centered on 50' Carrington longitude developed from
February to April and decayed afterwards. The former was observed by the earth-
orbiting spacecraft IHP 8 (Solar Geophysical Data, NOAA, 1976), but was hardlydetected bv the IPS observations of 3C144 (ecliptic latitude
= -1.4 ). This shows
that the edge of the stream was very sharp in latitude, as pointed nut by Sime andRickett (1978).
In 1977 and 1978, many active regions appeared at nid-latitides and no stablerecurrent stream was noticed. Therefore, we have calculated the velocity distribu-
tions fit to the Lower envelope of the observations of pattern speed. The northpolar high-speed stream of 800 km/sec extended down to 50 latitude in 1977, while
it extended only to 80* in 1978 (Figures 7 and 8). For these years, the lowerenvelope seems to be slightly asymmetical with respect to the day when the radio
source was closest to the sun. This is accounted for if we assume that the polarhigh-speed region was displaced by a few degrees from the pole, or that it wasgradually shrinking. It is interesting to note that the recurrent depression ofthe pattern speed was observed in 1977. This may have been caused by a slow-speedstream at mid-latitudes associated with an active region.
4. SUMMARY ANi) CONCLUSION
We have derived the distributions of the solar wind velocity for 1974-1978
using the 3C48 observations, and have shown the existance -f the polar high-speed(800 km/sec) source region and its extension to the eqatorial plane. As shown inFi gre 9, the north polar high speed region extended from the pole down to 45 51)latitide for most of 1974-1977 (we used 3(298 observaitions for 1975, as shown inFigure 5, because we did not make observations in lanuary-hune 1975), while the
boundary of tile region moved up to 8DO In 1978. The observations in 1979 showthat this region almost disappeared. The shrinking of the polar stream (faster
than 500 km/see), coinciding with a corresponding contraction of the polar coronalhole, h;s also been foind by Rickett and Coles (1979). It is well known that the
solar magnetic dipole field reverses near the solar maximum. Our results showthat the polar high-speed region shrunk without tilting from the rotation axis.This suggests that the dipole field does not rotate to reverse, but is apparently
vanishing. Tile behavior of the polar coronal hole is interesting, and it isimportant to continIii, monitoring,, the polar high-speed stream for the coming years.
Ill tlhis paper, we hav used the ohervations of only one source. If we add
the observa tlons fro,, two other sources, an ecliptic source and a southern hiigh-latitude source, we wi 11 he able to derive a more detailed distribution for both
160
D .97 ,
Ig? 9771978
Figure 9. Variations of the north polar high-speedregion during 1974-1978.
hei spheres. We are interested in north-south asyrimetry of the distribiaton4, :.iested by the UCSD observations for 3(161 in 1974, and we are planning to qtuldy
Acknowledgment
ic ire gritefl to Professor L..A. Coles for providing [PS dita i t i: SI.
References
- oCes, L.., and Rickett, B.I.: 1976, "IPS ohservations of the solar wind speed
out of the erlIptlc", 7. Gevoilys. Aes., 81, 4797.
olos, W.,\., larmon, .I.K., Lazariis, A.J., and Sullivan, .. D.: 1978, "Comparison of7'4 '.1: intorpanotary scintillation and IP' 7 observations of the solar windfiurina 1973", .1. Geoi hs. Res., J3, 3337.
1.,mi eu-,, P.A. , ind jew i h, A.: 196 7, "the solar wind out;ide the plane of theorl ti¢", 'fipt tI, .213, 343.
lwish, A., and Symonds, '1.1).: 1969, "Radio lnvestig,,;ation of the solar plasna","1iunit r .j ce -ci., 17, 313.
,inse,, R.T., Hlansen S.F., and Sawyer, C.: 1976, "Long lived coronal qtruictilr,.s
and rcurrent goeomagnetic pattern, in 1974", ?laneta ry -;lace .Sci., 24, 381.
lmiininor, :.: 1977, "Scintillation observations of the intprplanetary plasa"c, in'1.\. Shi.,, I).Y. Snart and S.T. Wit (eds.), Study of Travelling lnterplanetaryl'hennena/ 1977, I). Re idri, l)ordruht, Ti. 119.
,K in ima, T., I4,, lrui, !I., and Kojima, '1.: 1973, "On tlte analysis of the h o) ervi-tin; int,,rplanetary scintillation; obtained with three spaced receivers",i'r.!. Itroin. ,C., apa,, .25, 271.
161
References
Kakintiria, T.: 1977, "Observations of interplanetary scintillation: solar windvelocity measurements", in H.A. Shea, D.F. Smart and S.T. Wu (eds.), Study ofInterplanetary Phenomena/ 1977, D. Reidel, Dordrecht, p. 101.
Kakinuma, T., Washimi, II., and Kojima, ".: 1979, "IPS observations of the solarwind", paper presented at the International Workshop on Selected Topics ofTiagnetospheric Physics, Tokyo, March 1979.
Kojina, M., jatanabe, T., Ishida, I., laruyama, K., and Kakinuma, T.: 1979, "Afull-automatic system for the observation of interplanetary scintillation,
Proc. Res. Inst. Atmosph. Nagoya Univ., 26, 95.
Readhead, A.C.S., and Hewish, A.: 1974, "Fine structure in radio sources at 31.5Hliz-HlI", Mem. Roy. Astron. Soc., 78, 1.
Rickett, B.J., and Coles, W.A.: 1979, "Solar cycle changes in high latitude solarwind", paper presented at the NASA Symposium to Study the Solar Cycle andDynanics Mission, W]ellesley, Massachusetts, June 1979.
Sime, D.C., and Rickett, B.J.: 1978, "The latitude and longitude structure of thesolar wind speed from IPS observations", J. Geophys. Res., 79, 3841.
Vitkevltch, V.V., and Vlasov, I.V. : 1972, "Characteristics of interplanetary
electron irregularities according to observations in 1967-199", SovietAstron., 16, 480.
Watanabe, T., Shibasaki, K., and Kakinuna, T.: 1974, "Latitudinal distribution ofsolar wind velocity and its relation to solar EUV corona", J. Geophys. Res.7Q, 3341.
Discussion
Schwenn: In the ST [P-77 Proccedings we showed in situ measurents from both;elios solar probs during early 197t. There is clear evidence for the occur-rence of a sharp lat itudinal boundary confirming your model assumption. Did you
try to include real stream profile measurements into your model calculations?Kakinuma: Our model is very simple. ti should include a real strean profile to
obtain a more detailed velocity distribution.
"/
162
MCC
13. Interplanetary Scintillations at103 MHz
S. K. Alurkar and R. V. BhonslePhysical Research Laboratory
Ahmedabad-380009, India
Abstract
With a view to studying the structure and dynamics of the interplanetary medium,a three-station interplanetary scintillation (IPS) observatory is being developedin India. Operating with only half the aperture (-2500 m
2) at a frequency of
103 MHz, the first IPS telescope at Thaltej near Ahmedabad is making daily observa-tions of 8-10 radio sources. Parameters of the intensity pattern, such asscintillation index, autocorrelation, temporal power spectrum, moments of intensityetc., are computed. With a two-parameter curve-fitting method it is found that theIPS power spectra are best matched with an exponential law.
163
I. INTROI)ICTION
Observations of tile interplanetary scintillation of compact (Q") radio sourceshave been successfullV used to delineate the microscale (-100 kn) structure anddynamics of tile interplanetary medium in the ecliptic plane as well as outsideit. Such observations are also useful for deriving the structure of very compactradio sources, particularly at low frequencies, where the VLBI technique is imprac-ticable (Hewish et al., 1964; Cohen et al., 1967; Armstrong and Coles, 1972;Kakinuma and Watanabe, 197o ). In this note we present preliminarv IPS observationsmade at Tha tej near Ahmedabad.
2. OBSER\ ATIONS AN) ANAISIS
The radio telescope used for these observations comprises a filled-aperturedipole antenna array having an aperture of -5000 m" at 101 Mliz. (Only half theaperture was used for the present observatinns.) The r. . signals, after beingamplified through low-noise preamplifiers, are combined in a beam-rrming matrixyielding a 32-beam pattern. Tile size of each beam is 8' -W - 2' N-S. A selectedbeam is connected to a total-power receiver with a bandwidth of 2 MHz. An analoguerecording of intensity fluctuations of a radio source is made with a time constantof 0.1 s. After A/P conversion at a sampling frequency of 21 lz the data arerecorded on a digital magneti, tape and then prrolcssed onj in BM 600/44 computer.Normally about 20 min of useful data per day per source is availabl. So tarobservations on KC -48, 3C 144, C 147, 3C 161, 3C 196, C 237, 3C 273, 3C 298 and3C 459 have been made.
Autocorrelation and intensity power spectra were computed using an FFTalgorithm from successive 50)-s segments of data. After editing, the selectedspectra were summed to get an average spectrum. Usually eaci source was observedfor about 30 min, followed by an off-source recording of about 5 min.
In Figure 1 are shown typical IPS power spectra " r five scintillators. Off-source spectra' were taken into account in choosing tLe high-frequeny cut-off oftile on-source spectra (-3 liz and -2 iz for strong and weak scintillators respec-tively). The spectra were best described by an exponential law in the frequencyrange 0.3-3 liz for strong and 0.3-1.5 Hz for weak scintillating sources, withindices of 2.6 for 310 elongation and about 4.0 for elongations exceeding A0'.
The indices are plotted versus elongation in Figure 2 for four sources. Atelongations around 200, 3C 48 has an averag, index of 2-3, which increases to about4 at higher elongations beyold about 200 and is scattered around an average of 4.
The first and (square-root) second moments (fI and t-) of ie intensity spectrawere computed using the relations
fl = j fc f P(f)df / f P(fdf
and I-1 = If NI f)d 1 jf l I)df, 1
164
IIIIII ... II I I N .>
S4 L . C. 141 aC4
3-6-79 3- 4-1-5-79 -5-79
I.
o4 . • , 0 1 , _ ._. . -. .
FRElIQUENCY INZI
Figure I. IPS spec..tra for five source.s. The thick solid lineis an exponential fit for 3C 48 data over 0.3-3 Hz and for
•3C 48
S3C 196
Z6 0 3C 237wA 0
tN;
4 4 0
0.- . .L.
44IIiI
3C 4.
0 i0 20 30 40 s l
ELONGATION (DEG)
Figure 2. i dices f exponential fit vf. e4vngation.
othersoures ovr 0.-.5H
where 1(f) is the power spectrum estimated tip to the cut-off frequency fK at whichtile spectrum meets the off-source noise spectrum; fl and f: are measures of thewidth of the spectrum, the former being less affected than tile latter by the noisein the tail of the spectrum. Further, 11 gives more weight to large-scale com-
ponents and is therefore more accurate because of the increased signal-to-noiseratio (Readhead et al., 1978). For the IPS data taken at Ahmedabad f, and f, werearound 0.8 and 0.9 Hz, and these values remained nearly constant for elongationshigher than 30'. The quality of the data can be judged from the ratio f21f1(-l.13), which in the case of a noise-free exponential spectrum is equal to "2.
:1. I)ISCUISSION
For the 103 11,0tz IPS data taken at Ahmedabad for solar elongations higher than20' the power spectrum index is 2-4 for strong and about 4 for weak scintillators,the former being comparable with tile results of Milne (1976). More precise deter-mination of the background noise spectrum is necessary for accurately fixing thecut-off frequency of the scintillation spectra.
Acknowledgments
We express our sincere thanks to Professors D. Lal and S.P. Pandya for showingkeen interest in this work. We are grateful to all our colleagues in the IPS groupfor their dedicated efforts in developing the IPS station and also for providingcomputer programs for data analysis. Financial support for this project came fromthe Departments of Space and Science and Technology, Government of India.
References
Armstrong, J.W., and Coles, W.A.: 1972, "Analysis of three-station interplanetary
scintillation", J. Geophys. Res., 77, 4602.
Cohen, M.H., Gundermann, E.J., Hardebeck, H.E., and Sharp, L.E.: 1967, "Inter-planetary scintillations. II. Observations", Astrophys. J., 147, 449.
Hewish, A., Scott, P.F., and Wills, D.: 1964, "Interplanetary scintillation ofsmall diameter radio sources", Nature, 203, 1214.
Kakinuma, T., and Watanabe, T.: 1976, "Interplanetary .cintillation of radio
sources during August 1972", Space Sci. Rev., 19, 6!1.
Milne, R.G.: 1976, "Interplanetary scintillation power spectra", Aust. J. Phys.,29, 201.
Readhead, A.C.S., Kemp, M.C., and Hewish, A.: 1978, "The spectrum of small-scaledensity fluctuations in the solar wind", Mon. Not. R. Astron. Soc., 185, 207.
16b
00
14. Analysis of IPS Data Using APeriodic Phase Screen Model
for the Solar Wind
A. Basu and S. K. AlurkarPiysical Research Laboratory
Ahmedabad-380009, India
Abstract
The intensity pattern obtained by diffraction from a periodic phase screen offinite dimensions nas two distinct length scales: the width of the diffractionmaxima and their interpeak distance. These can be uniquely related to two lengthscales in the screen, namely the lateral extent of the screen and the phase period.
In the weak scattering limit the autocorrelation and the Fourier transform ofsuch a pattern are dominated by terms depending on the total size of the screen.We justifv our assertion that the interplanetary medium (IPM) acts like a screen offinite lateral extent, and show that in this case the second moment of the powerspectrum does not yield the desired information on irregularity sizes in the solarwind. This must he extracted from high-order terms, in particular from the dips inthe power spectrum.
The results of the present model and Salpeter's (1967) random phase screenmodel are compared with observations. The validity of the model as an idealization
167
of tile 1PM is also examined.
Ilhe valuvs at irregularity size calculated from ohservati ,ns taken atAhmed;l Iad are I ltow hundred kilometres on tle basis ot this model.
I. INTRODUCTION
Radio waves from astronomical sources reach the observer on Earth afterdiffraction through a translating screen oi varying refractive index formed by thesolar wind. In the weir' scattering limit the scale of fluctuations in the mediumis usually determined from
(i) the spatial auto-correlation length of the intensity pattern, given the solarwind velocity, or
(ii) q- 1, where q is the first moment (Readhead et al., 1978) or square root of
the second moment of the power spectrum.
The power spectrum or inverse Fourier transform of the intensity F-I) gives tilepower distribution in tile fluctuations over a range (f waVe numbers (Salpeter, 1967).
We will show that if the medium is finite in lateral extent, the major contri-bution to the above quantities comes from !., the screen size, rather than -, thescale of irregularities in tile medium, assuming the SCren to he phase periodicwith I, tile amplitude of phase variations, about unity.
The necessity of considering a finite screen arises hel',cause the elfective sizeof the screen is always limited by several factors.
(i) For very weak scattering . = Ls = 2z- 5s, where ", is the angle of scattering.
(ii) If the source is not monochromatic, tie temporatl cohrence, of the source, or
finite bandwidth of ti reo' iver, sets a limit to 1. beyond which coherence is de-stroyed. We assume tit the handvidth can he made softiciOntlv snall so that thislimit need not be considered.
(iii) For all incoherent source (e.g. a star) the lateral extent of coherence oftile wavefront at aI distance z from the source is (Champencv. 1973)
L, = 1.22',z/D) 1.22 /,
where % is tilie anglIar size of the source and 1) is the diameter ot the source. If,t is sufficiently small so that Lc l.,, the source scintillates. Large sourcesare non-scintillating, since incoherent contributions destroy scintillation when
Lc .
If L becomes comparable to Z, then the assumption of a "random" distributionof irregularities is open to question because It is valid only in the presence of
a large number of irregularities within the diffracting region, i.e. when i ' I..
If tile spectrum of Q is truncated by I., and in addition the power hierarchy ofthe irregularities is such that the largest sizes contain the maximum power, thenthe number of irregularities having the highest power within the scattering region
I (i8
may be very small. In such a case the periodic phase approximation may be closerto the actual situation than the random phase approximation.
In the following we calculate the effects of diffraction by a periodic screen.
Only the Fraunhofer results are presented (although this imposes the ratherstringent condition L
2/4z < A/20 (or L z 3 x 102 km for z 1 A.U., - 3 m) because
of the extremely simple linear relations that emerge between pattern scales andcharacteristic scales in the medium.
2. DIFFRACTION FROM A THIN ONE-DIMENSIONAL SCREEN
When a plane wave of unit amplitude is diffracted by the infinite screenf(xO), the complex amplitude is (Champeney, 1973)
f(X,z) = + ei2 /A e-i2v (Xi z)e f(x,O)dx (1)
and intensity
I(X,z) = 1/Az Fx/ z[f(z ,0)j , (2)
where A = wavelength of diffracted radiation,z = distance of observer from the screen,
x,X = coordinates at the screen and observer respectively,
FF- = direct and inverse Fourier transforms,X/Iz = Fourier conjugate of x.
The inverse Fourier transform of I(X,z) is
r' 2TriS X F=- w, c2rS EX,) [f(x,0)]J dX
or
"2,-ir(W/z)/ X ~~,)= F 2/i I d(X/Az) (3)
where AzS has been replaced by r. By the autocorrelation theorem (Bracewell, 1965)the right-hand side of Eq. (3) is the autocorrelation of f(x,O).
F/ Af( 50 ) (r)
<f(x,O)f(x+ r,O)> (4)
169
Equation (4) suggests an easy way to determine the autocorrelation of an arbitraryinfinite screen if the spatial intensity distribution in the Fraunhofer region isk nown .
If the screen is of finite lateral extent, the new screen function
f'(x,O) = f(x,O)W(x/L), where W(x/L) is an envelope function of width 1. that tendsto zero for large values of x. The intensity is the convolution
2
l(X,z) = Il,\z IX/,zf(xO)1 * FX, z W(xIL)] (5)
and
Fr/A Il(X, z)J Af,(x.o)(r)
<f (x,O)W(x/l) f(x + r,0)W(x + r/L)> (6)
It can be seen from Eq. (5) that thec intensity pattern is the convolution oftwo independent functions with different scales. Since the scale of a function isinversely related to that of its Fourier transform, the first term on the right-hand side of Eq. (5) gives a pattern scale proportional to .z/. and the second terma pattern scale \z/l.. In other words, if there are two independent scales in theintensity pattern, the shorter is inversely propor-ional to the overall size of thescreen while the longer is inversely proportional to the irregularity scale of thescreen.
2.1 Diffraction I- a Thin One-Diniensional Periodic Phase Screen
I1 f(x,0) = ei. cos (2 ix/ .* ,lnd
W4(x/i,) e x
(7)= 0 x, >"TI
the expli, it form oI the intensity is determined by Ratcl if-, (145f) for in infiniLt,screen. For I i ni t- 1.,
n=--
As shown in Figure 1, this intensity exhibits a scries of diffraction maxima ofwidth 21z/I. and separation z/l.. In(") goes rapidI, to zero for n and thissets a limit to Lhe order of diffraction maxima actually seen.
170
n=O
I (Xz)
XZ,,j n=1
nz2
2XZ /L
Figure 1. Diffracted intensity L(x,z) at a distance z from a phasescreen f(x,O) = ei
1 CO 27X/;Z truncated at X= 41 2 1_ for =I
2.2 The Inverse Fourier Traiisforiii of I and( the Second Nfolneiit of F -(1) for a Periodic Phase Screen
From q . (8):
Fu = ( L J1(:)A( SIZ) e 2i (n'-USI)
or
where -r/I.) II-r Ii<-r1,
=0 rI> 1
AWLII) is nothing but the autocorrelation of W(x/L) as expected fromi Eq. (6).
Fo r I,
171
(see Figure 2), and the normalized second moment of FS(1)
S = 6 T6 (II
From Eqs. (10) and (11) it can be seen that the dominant contributions to F-I andS, come from terms independent of for * 1.
t
Re Fs( I)
SN
L 6
L
r=L
r XzS
Figure 2. Fourier transtorm (, I(X,z) tor 1.
= -!r=:r r'
FS II = 0
2.3 The Aistocorrelation Fijwltion of Intensity. for a Periodic Screen
Simiai r I y w . Iind
172
A1 R <1('X,z)1(X +R,z)> ~ ~ C A X
nII I , n, X12 (XI +A) 2
(12)
J J~ siw X' sin2(XI +A, dX'
n, n~
wilcr" X, I'.
1 i R
I t can be shown t ha t
sinl N sinll + sill 2.dx x +
I J)(;n
in* s rw torm A ~r ,,s po nd ing L n n 0) is i ndisndon - I
I sinAoo I -
LR
[II, '"Itspoiids to th' IieontriI I lohi ,Ith auocorre Iat ion funct ion and hasos-il liojs on .1S alto 'z/I At thO tiist MiniMUM, and 1, = -7/K. Ihisrokit ion is used to determinc 1. from A (K).
lki iii tortli q) to h, ncx-t order Ior I,
17 3
-! 1 - -
2 + .. . ... I- -23: ii 2 <
.1 , ".,I\ t L
Ai (R)
XZ1
Fi.ur, . ,tr -'n t ion ,f int ,sit ,
I I '' ( :
A ' .....[ Ill)
ii + 2=19
W~lL' I
2-.,
174
3. DETERMINATION OF PARAMETERS OF TIlE SCATTERING SCREEN FROM TILEOBSERVA.TIONS
:3. Dieeriinatit i of i' and r 'or a Periodic Screen
From Eq. (10) it can be seen that the minima in Re F/(1) are situated at
. etc. (18).Z 2z ' 2z ' 2z
That is, the distance between the dips in F-(i) gives a direct measure of Z,provided 1 - . If * 11, additional dips are obtained at intervals Lr = Z/n,which would make the determination of k difficult (n - ¢ - 1; 2 < 1 < 10). : can
be obtained from a set of relations for each dip in F-(i).
- ,7/"z I)!r k / ( 2J (¢) (
-_ = 3, 5... (19)F r=0 2, ) 2
The width of F- (1) corresponds to L.
In the autocorrelation function the distance to the first sidelobe is \z/Z.
Although * can be determined from AI(R) in principle, this method is likely to beunsatisfactory because the autocorrelation is normalized by its value at R = 0,which has a large contribution from random noise.
3.2 )istribution of Power in )ifferent S'ales for an Aperiodic Screen
The distribution of scale sizes 9 is estimated from the frequency distributionof the spacing betwecn successive peaks in the scintillation pattern. If the time
lag between peaks is :, and N(-) the frequency distribution of i, the power -7(q)in the phase fluctuations as a function of wave number is proportional to N-(7),where q = 2"7v"/ z. Figure 4(a) shows a sample histogram of N(T) vs. I and Figure 4(b)
givcs log N (7) as a function of log -. The steep fall in the distribution for lowvalues if - is due to the finite size of the screen, which filters out irregulari-ties larger than L.
. (OIPARISON OF' Q1 ALITATIVE FEATURES OF TIlE MODEL WIT!I OI.SERVEI) FEXTIRESOF II'S
Ih model for the solar wind is
f = i ( ;)(-)s(2
t: x/;)
t ( X,
It is assumed that (M) is an increasing function of 2 for 2 < L. The specificform of l(q =2 T/I) can be determined experimentally from N(O) vs. i (Figure 4).
So- logN 2 )
so (a) 4 Cb)
70\
O60 3 K
N(t) 50 -
40 -2
30 x X
20
X10
I I [ l l , I - I
2 4 6 a 1o 120 0 2 0 4 0 0. 1.0 I-2
T - Sec Iog 1 -
Figure 4. (a) Histogram of N(T) vs. r. The wave namber q equals2nv/z. (b) Log N2(T) vs. log T. The corresponding range of q is
1.4 , 10- 2
km-1
to 8.7 ' 10-2 km
-l
The size of the screen L = L= 2z<s, where
= r n e 2 (Scheuer, 1968),
r = classical electron radius,
An = density fluctuation of elecLrons,e
I) depth of scattering region.
The largest irregularity within the scattering region is x, 1,. and since Z()increases with k, it is also the dominant scale.
1.1 Oliersed Feahreq of IPS
(i) From tlhe mode I w, find that the depth of fading varies proport Iona II y withthe inverse of the width of the diffraction peaks, and thus it varies proport ional lyto k. We also tind that the scintillation rate is proportional to the inverse of
176
I-
the distance between peaks, which is proportional to
/ z Ls/Az
as observed in the case of scintillation.
(ii) Again from the model we can predict the variation of scintillation index m
with elongation t-. We expect that with decreasing elongation Ss would increase,and consequently m also would increase (Eq. 8) until L, = Lc . Beyond this, an..
further increase in L, would cause m to decrease owing to incoherent contributions
from the edges of the wavefront.
This is in qualitative agreement with the variation of m with c as seen inFigure 5.
EX 0.4-UJ
Z0 -0.3
_
Z 0.2z- O,2
U)
0 20 40 60 80 100 120
ELONGATION c (deg)
Figure 5. Variation of svintillation index with solar elongation
The souzrce size xcan he determined from the point at which m = Nax 'Yhewidth oif the peaks is equal to iz/l 0 , which is approximately equal to 2'.,11.22. Interms of the cut-off frequency V = v/Ic
=1.22v m aat =
177
NONNI"
This is in agreement with Cohen et al. (1966) except that according to our argumentsthe source size can be obtained only from the point m = mmax. Elsewhere the cut-off frequency is determined by L. and is unrelated to the size of the source.
(iii) The maximum of m vs. c occurs at smaller elongations for shorter wavelengthsand smaller source size (Figure 5). At m
= mmax,
L = 1.22A/a = z(D/0)1 r An A2
s e e
If a is assumed constant, then a shorter A requires a higher Ane, which mean, thatthe maximum of the m vs. c curve shifts toward smaller c for decreases in .\. If Xincreases, m = mmax shifts towards larger elongations. It remains to be seenwhether quantitative agreement can also be obtained with observations.
(iv) According to Salpeter (1967) the dips in the power spectrum are due to asin q
2 factor (the Fresnel filter). The spacing between the dips is then expected
to decrease with increasing q. In the Fraunhofer region the dips are expected to beequally spaced, as we generally observe. Buckley (1975) and Lovelace et al. (1970)also mention that the dips in the power spectrum do not agree with the theoreticalvalues from Salpeter. This may be due to the finite size of the screen, whichmakes the Fraunhofer approximation adequate except at very large elongations.
5. RESULTS
Some sample values of Z and L calculated from IPS observations taken at UCSD
using the results of our model are shown in Table 1.
Table i
K L/A z
(km) (km) (km) (km)
20-30* 1.8 ' 102 3.4 10' 1.53 × 107 1.36 , 108
40 1.33 10" 8.85 102 2.94 106 1.15 l0
60 1.31 10-' 9.17 10" 3.00 10" 6.90 × 107
70 1.3 - 102 4.t7 10" 1.36 10" 5.57 - 107
96 " 1.5 1(o 7.8 101 2.93 l ]i' 1.5 107
*The true ir:,-gularity scale is -( -I)Q. if 2 10.
tFraunhofer approximation and localized thin-screen
approximation not expected to hold.
178
At small elongations, if P ) 1, then the true irregularity scale is ( -1)l.
Taking this into account we see that tile largest irregularities that are observedare of the order of hundreds of kilometres.
In Figure 4(a) the histogram (c - 40*), prepared from IPS observations taken
at Ahmedabad, gives the distribution of (Z) over values of . ranging from -10 to
500 km. There is very little power left in irregularities below li00 km. The peakpower is at 400 km, above which truncation effects set in.
After the completion of this paper an earlier work (Briggs, 1961) was brought
to our attention. The conclusion of Briggs was that the scale measured from thepattern autocorrelation is either 2 or L, depending on whether we have z 4 1, orz > XL/k. Since our calculations are in the Fraunhofer region they appear to be in
agreement with Briggs.
Using this criterion we see in Table I that for IPS we are usually in the"far" region as defined by Briggs. Our results may therefore be of relevance to
IPS observations, even though Fraunhofer conditions may not hold strictly.
0. SUIMARY
Ii) The auto-correlation function of a diffracting screen and its limitinenvelope can be obtained from the relation
AfkxO)W(x/1)(r) = Fr/ z(
kii) In the Fraunhofer region, pattern scales are inversely proportional to thescales in the medium. The shorter scale is proportional to L- and the I cnkerscale to ;71. To determine , the longer scale must be selected.
(iii) Since tile autocorrelation function and second moment of power spectrum pre-ferentLially pick out the smaller scale L
- 1, they may give erroneous 'stimates Of
(iv) Source sizes can be determined from the point m = mma x
Angular diameter -.. (width (of diffraction peaks)z
1.22vf
where f is the frequency cut-off in F-(.C
(v) From Fr/;z(I), the value of can be determined by taking the ratio of
at each minimum with Its value at r = 0, if 0 2, and using the relation
.) r .k.. . 1- k = 1, 5 ...
F- 0y ! j 2LI
179
(vi) If 1 " 1, the dominant Jn occurs at n I- 1 for 4, 5 and for = 10,JI and J are comparable.
For b < 10, the true irregularity scale will be (i-l)Z, where Z is measuredfrom F-(l)
(vii) The model cannot be used if the power in several scale sizes is comparable.
Acknowledgments
We are grateful to Professor W.A. Coles for the use of some IPS data taken atUCSD, and to Dr. B.H. Briggs for comments and for directing our attention to hispaper.
References
Bracewell, R.N.: 1965, The Fourier Transform and its Applications, Mcraw Hill,New York.
Briggs, B.H.: 1961, "Diffraction by an irregular screen of limited extent",Proc. Phys. Soc., LXXVII, 305.
Buckley, R.: 1975, "Diffraction by a random phase-changing screen: A numericalexperiment", J. Atmos. Terr. Phys., 37, 1431.
Champeney, D.C.: 1973, Fourier Transforms and their Physical Applications,Academic Press, London.
Cohen, M.H., (;undermann, E..l., tardebeck, H.E., Harris, D.E., Salp.ter, E.F., andSharp, L.E.: 1966, "Radio sources: Angular size from scintillation studies",Science, 153, 745.
l.ovelace, R.E., Salpeter, E.E., Sharp, L.E., and larris, D.E.: 1970, "Analysis ofobservati'ns of interplanetary scintillations", As trophys. .1., 159, 1047.
Ratcliffe, I.A.: 1956, "Some aspects of diffraction theory and their applicationto the ionosphere", Rep. Prog. Phys., 19, 188.
Readhead, A.C.S., Kemp, M.C., and Hlewish, A.: 1978, "The spectrum of small-sci ledensity fluctuations in the solar wind", Mon. Not. R. Astron. Soc., 185, 207.
Salpeter, E.E.: 1967, "Interplanetary scintillations. I. Theory", Astrophys. J.,147, 433.
Scheuer, P.A.G.: 1968, "Amplitude variations in pulsed radio sources", Nature,218, 920.
180
MO
Discussion
Zlotnik: Could you say anything about the shape of irregularities? Are they
iostropic or elongated? What about the scale sizes of irregularities?Alurkar: They are elongated, with typical ellipiticity ratio of around 1.8.
IPS determines scale sizes in the range 10 to 1000 km.
181
15. STIP VII Observations ofthe Northern Solar Corona - A
First Report on the Scope ofData and the Early Results
Thomas A. CroftSRI International
Menlo Park, California 94025, U.S.A.
Abstract
During late August arid early September, 1979, four interplanetaryspacecraft passed through superior conjunction. Within a short time theirradio signals all passed through the northern corona of the sun on the wayto earth. Because this event was foreseen, the months of August and Septemberhad been declared to be STIP Interval VII so that observers, using earth-boundinstruments and instruments aboard other spacecraft, could acquire data to becorrelated with measurements derived from the four interplanetary spacecraftsignals. These conjunctions offered an excellent opportunityo observe thepl. ma within a few radii of the sun, at latitudes between 20 N and the Pole.
In this first report on STIP Interval VII, we stress the type and extentof the measurements that were obtained from the spacecraft. Four types ofsignal processing will yield different types of information about the solarplasma through its influence on the passing radio signals. These types ofprocessing and their resulting data are described. A severe limitation on
1d3
the schedule was the high percentage of available tracking time devoted toPioneer 11 its it passed Saturn on September 1, 1979. For this purpose, as thehighest priority, observing time was assigned to NASA's limited number ofreceivers. Coverage of the other three spacecraft was achieved through the
hard work and enthusiasm of personnel associated tith NASA's Deep Space Net-work (DSN) of receiving stations.
In this report it is not possible to include many of the experimentalresults because the processing is only partially completed. We do provide acomprehensive guide to the expected results so that the scientific communitymay have a schedule of our observations.
1. BACKGROUND
As a radio wave travels in interplanetary space, its velocity is affectedby the number density of the free electrons. Because the solar wind ismostly ionized hydrogen, it follows that plasma density is almost equal to theelectron number density. (A correction can be made for helium to increase theaccuracy by a few percent.) The solar wind and corona influence passingradio waves in a variety of measurable ways, most of which are manifestationsof changes in the phase velocity, i.e., the speed of the wave fronts.
With American spacecraft, the signals to earth are S-band (2295 Mliz) orX-band (8415 MHz). At these frequencies the waves travel slightly faster
than the speed of light by an amount that is, for all practical purposes,proportional to the electron number density. Signals that are impressed uponthe radio waves by some form of modulation must travel at less than the speedof light, i.e., at the group velocity. Both phase and group velocitiesdepend upon the frequency of the radio wave; plasma effects are much "rger
at the lower frequencies. Other radio propagation phenomena, such as
absorption and scattering, play no part in the program of measurement andanalysis we are undertaking.
1.1 Observalde Effers
In order to determine interplanetary group velocity, we can impress thesame modulation simultaneously on X-band and S-band transmissions from the
spacecraft and then measu- 2 their relative arrival times at earth. The delaydifference shows how much greater the effect was at S-band than at X-band.The delay difference between the two bands is attributable to plasma. Becausethe group velocity is decreased below the speed of light by a small quantity,which is again proportional to the electron number density, the delaydifference can be converted into a measurement of the electron content alongthe line of sight. (The delay is an integral of delays all along the path.From this we derive an integral of electron number densiLy all along the path,
184
which is the electron content.)
The measurement of differential group velocity is disadvantageous intwo ways: (1) It requires the provision of expensive special equipment onthe spacecraft that is often not used for any other purpose, and (2) it isinaccurate in comparison to the corresponding measurement of differentialphase delay. Unfortunately, it is not possible to measure the phase delay;we cannot measure the relative time of arrival of waves at the two frequenciesbecause any distinguishing characteristic imparted to a wave that makes up thesignal must travel at the group velocity, by definition. Instead we mustmeasure the frequencies of the two signals, that is, the rate of arrival ofwave fronts. If the electron content increases, then the frequencies willchange accordingly. From frequency measurements we can determine the timerate of change of electron content along the path. This can be integratedwith respect to time, yielding the electron content, but the constant of
integration can never be determined except by independent measurement of thegroup delay.
The differential phase delay offers advantages of equipment simplicityand greater accuracy. The best results are achieved when both systems areimplemented, as has been done in a solar-wind experiment flown aboardPioneers 6 through 9 (Croft, 1979).
Small-scale turbulence and irregularity in the corona (or solar wind)causes correspondingly small-scale distortion in the wave fronts that traversethe medium. The resulting interplanetary scintillation (IPS) is most oftenmeasured by the signals from radio stars, but spacecraft signals provide amore readily interpreted result because the source of the signal is known tobe vanishingly small (in areal extent) and temporally stable.
1.2 Cl.sed-Leop and Open-1,., p Records
During the reception of spacecraf, signals by the DSN, controllableradio parameters are optimized for the detection of signal characteristicsthat maximize the likelihood of telemetry success. A fairly crude measurementof scintillition may be obtained from the "closed-loop" records. (This namerefers to the fact that the phase-lock loop is fully opLerational at the timLof signal detection.) Such noise data may be available during all the hoursthat the receiver operates. 'lhe great virtue of such records is that th y canbh obtained for long, uninterrupted intervals at modest cost.
At times of high priority, special recordings are made by highly linearreceiving apparatus in which no effort is made to detect the signal. Insttad,the signal plus all the noise in predetermined bandwidth is carefully filteredand the resulting noisy waveform is then converttd to digits and recordeddigitally. By this means, the scientist can gain access to almost anyparameter associated with the radio signal with almost no degradation. Thebest scintillation data are obtained in nis manner, which is called "open-loop" recording because there is no phase-lock loop in action at the time of
reception. The practical limitation lies in the time and expense ofprogramming and operating a computor to simulate the action of a radio re-ceie, r, dt tect the signal and process it t-' extract the parameter of interest.
185
In this report we call the lat ter type of data a 'linear' rcciordl iug aniddesignate it with symbol L. Those recordings of till- d-et c e sif!nal1 1biili Itthle time of actual reception by closed- loop rt-ct i vers art. givt-n tit, syiibo D).
1.3 Tile Mean andi V'ariance of I'lasizia IDensili
In another, broader , view the spac-craift int. sorin-it s ar, s, , it tof L IIinto two classes: (1I) Some of the measurement s re fat, t o Lh fit'(I gre oft localIirregularity in thle plasma ; thereforer, we consider Lt.t theseR m, isor, itent I.provide information about the variance of thle densitv. In 11: ilikt mannierof thinking, the differential phase and group del iv obs, rv~ . iiis preY id'evidence about the mean value of the denlsity.
2. SPACECRAFT ALIGNMENTS IN .X XS NI) ETIIh
The usefulness of these spacecraft lies in the fac.t that their iadjopath to earth passes near the sun. As is suggested by the conceptual sketchin Figure 1, the signals depicted are sensitive to the, solir plasm-.a. Thelirgreatest sensitivity by far is in tile vicinity of the sunl. Till tendeonly forthe plasma density to vary roughly as the inverse squl.- of tli, di ctaiiee fromthe sun creates a natural weighting of th0, kiu's spla.smi~ilu Lvtewiid~ thelraypath perihelion. As a consequenIce, Oute of till nest ifilperteait . meltparameters is thle perihelion distance, perhaps tasilv vistializ-d inl t. ms ofthe Son-Earth- Probe (SPEP) angle or the( appa rent d is t,1ne eo1 till SpaCLcraftfrom the stil as thev bothl appear in the s'Ky to the imaginat ire ohs..rv.-r. Auseful graphiical IpreSenltation is thlen thet Vv ec of ti., ..pi- craft as seen f remearth, drawn in a cross sect ion perpendicular to tile line of sighit and 0nlltaining thle son. We will use such a drawing, excniplified hrv Figure 2, whichillustrates tile closest approach from amiong tie, four paths. III this pict are,the spacecraft appears to move from left to right icith t he paLSsage1 Of time',but the motion is only relative to thle sun. If thet ~iettire Wete draiwn in aninertial frame, then thle Voyager woulId more s lowIV vLo tile left, but tile Sunwould move more rap idly to the left. These msot ions arte tan scd hr- the comlb i nedmovement of the Voyager and of thle observer tin earthi; becauise till angulIar rat.of mot ion of thle earth exceeds that of thle VerAger,~ thet latter a- PPearS tomove to thle right re lative to the stil. The same no Ids true tO P'ioineer 1 I andVoyager I, bitt Piloneer V'enus is nearer thll Sun aiid thus has,. higher ailgo bIrrate than earthi. It appears to move to the Left. These act ioins are d iagraviried1in Figure 3, whichi depiictIs thel four traijectories for alImost three monthis. Thelpositions of lielieis I and 2 are shfown in Frame. (0e) Of the figure. InI thIL listframe is shown thle rielitive Pos it ion of all six spacecraft oIn one Part icularday at about thle midpoint of the interval of interest. It is importantL tovisualize that each of these spacecraft is inl Miot ion at a di fferint rate liil adifferent direction; all the directions are tiearlNy list-west.
w >Q)J
00z4
ui m 41
V)-
4> 4.
o~-4j0'
~U4 -r
187.
PATH OF VOYAGER 2 AS VIEWED FROM EARTH
TYPICAL DSN 4 HAfTRtACKCING PERIOD
TRACK OF VOYAGER 2 IN CORRECT RELATIVE SCALE TO SUN. PROMINENCE.ALSO 10 SCALE. TAKEN FROM SKY LAS ATM 3 EUV IMAGE OF SUN
Figure 2. The close approach of a signal to the sun is evident in Scaledrawings of the View fromi Earth (cxarnple provided by D. P. Homes).
r AD A130 687 PROCEEDINOS OF THE STIP (STSDY OF TRAVELLNG
NL PA NEAHNMN NEGOHSC HANSCOM AFB MA M V SHEA ET AL 27 DEC 82
UNCLASSIFED AFGL-TR82-0390 FG32 N*rnuuuuomiouuumEmhhmhhhEmhhI
III L *3.2 2.
I liii1.811111.25 .4II~ -I 6
MICROCOPY RESOLUTION TEST CHARTNA.OA DSCAU Of ST.hDonOS -,6- A
NORTH
JULY 30 AUGUST SEPTEMBER
10 20 30 10 20 30 10 20 30
20 10 SUN 10* 20o 3oO
(a) VOYAGER 2 ON TRAJECTORY JSX - MEASURES 5- AND X-BAND SCINTILLATIONAND ELECTRON CONTENT
JULY AUGUS SEPTEMBER
10 20 30 10 20 30 10 20 30
SUN
Ib) VOYAGER I ON TRAJECTORY JST - MEASURES S- AND X-BAND SCINTILLATION,AND ELECTRON CONTENT
JULY AUGUST SEPTEMBER30 10 20 30 1o-J- 20 30
SUN
Wc) PIONEER 11, PASSING NEAR SATURN ON SEPTEMBER 1- MEASURES S-BAND SCINTILLATION
4SEPTEMBER AUGUST JULY
SUN
IM4 PIONEER VENUS, ORBITING VENUS- MEASURES S- AND X-BAND SCINTILLATION.MIGHT MEASURE RATE OF CHANGE OF CONTENT.
HELLOS 2 HEUOS IJULY AUG SEPT JULY AUG SEPT
10 20 30 10 20310 10 20 30 102030 10
SUN
I) HELIOS I AND 2 - MEASURE S-BAND SCINTILLATION
OP11 PvJS T 0 40 ® J s x E LlO S 2E
,i I , , , I I , 1 , , 1 e 1 I I 1 4 1 I 1 1 1 I1' I30* 200 10 SUN 100 200 300
It) AN EXAMPLE - ALL SPACECRAFT ON AUGUST 20. 1 17
Figure 3. Positions of six spacecraft relative to the sun as seen from
Earth.
18()
2.1 Encounters Between Spacecraft Along ||eliwentric Radials
Because three spacecraft moved right and one moved hft, Lhere, wore ti hreoccasions on which two spacecraft radio paths lay along tht. same radial lintfrom the sun. Figure 4 illustrates the first of these occasions which, in asomewhat humorous vein, have been termed "close encounters". (EOS, 1978.)In this figure we see that the radio paths from Pioneer Venus and from Voyager2 were in the plane that contained the center of the sun. The plane wastilted 28.40 N of the ecliptic plane. The plasma which traveled throughthese radio signal paths is assumed to be traveling radially outward from tilsun at a solar latitude of about 280. The important point is that this sameplasma passes through both radio paths in turn. To my knowledg, no suchencounter haa been observed previously.
The value of an encounter is twofold: Firstly, it offers a means forstudying the comparative scintillation on the two paths; therefore it is allavenue for determining the evolution of turbulence in the corona while theplasma makes its transit through this important region. Secondly, we hope tofind some characteristic delay between the two records that will indicate thetime of travel for the plasma to make the 3.2 R trip between the two paths.If this is possible, then we will have an independent measurement of theplasma velocity that can be compared to the same velocity inferred from thescintillation by an established approach. The derivation of a velocity valuewith high confidence would itself be of value; however, the most useful pur-pose served would, perhaps, be the calibration of the scintillation approachfor deriving velocities.
2.2 The Exceptional Pairing
The closest encounter of all three was that between Voyager I andPioneer Venus depicted in Figure 5. The plasma transit between paths on this
day may have occurred in only 10 or 20 minutes. Til switch in DSN recordingbetween spacecraft had to be rapid, far faster than the ont. hour that isconventionally set qside for this task! I travelLed to the Goldstone Tracking
Station on August 22, 1979 to describe this quick-change objective. With thehelp of the enthusiastic staff of the station, a series of short cuts weredevised so that the transfer was made in only 50 seconds. Considering thecomplexity of the electronics and the enormous mass of the b4-me-ter dish that
had to be steered in such a short interval, the 50-second switchover is arecord that may stand for some time to come. It was only possible through
the skill and innovative approach of the operators.
2.3 The Pioneer Pairing
Another, more distant encounter, which is depicted in Figure b, involved
the two Pioneer spacecraft. This and the August 23, 1979 encounter suffered
from a weak Pioneer Venus signal. Its high-gain antenna had been disabledintentionally from August 20, 1979 onward so that the Pioneer operators
could devote their full attention to the Saturn Encounter. We look forward
with considerable interest to the processing of these signals; we hope that
they will prove strong enough to provide useful data.
190
SUN
Figure 4. The radial alignment of two signal paths on August 17, 1979.
191
-- =Radio Signals to Earth
OAR O 0
SUN 1
Figure S. The second radial alignment; the closest encounter of the three=
192
UN
z -
0.-
00 m.
('- a,
-,4
- -IC
A -r--
193
2. A Useful Guide
A summary of the spacecraft positioning during the key parts of STIP VII
is shown in Figure 7. This is drawn to scale, as it would appear to an
observer on earth if the spacecraft were visible and were photographed each
day at 00 UT. Those readers who are contemplating correlative data analysis
will find that this figure, together with the tables of the time coveragewhich are given later, gives a summary of the spacecraft records.
One objective of collaboration will surely be the tracking of solar active
regions wherein these spacecraft records offer evidence of the ejection of
high speed, turbulent, or dense plasma. The emphasis must be placed upon
those active regions at latitudes north of about 200. With this limitation,
the Voyagers and Pioneers show much promise.
3. ILLUSTRATION OF THE FORMS OF OBSERVATION DERIVED
From fluctuations in the frequency of the arriving signal, detected by
the closed-loop receivers, the DSN can calculate and record on magnetic tape
the "Doppler noise." Information about this parameter is given ky a series
of articles in the bimonthly "Deep Space Network Progress Report ," which is
published by the Jet Propulsion Laboratory of Pasadena, California. Anexample of the data obtained by Pioneer ii on September 4, 1979 is presented
in Figure 8. The lowest record shows Doppler noise with a sample averaging
time of 0.1 second. For the successively higher records these averaging timesare I second, 10 seconds, and I00 seconds. From a comparison of these, onecan derive the spectral index by a method illustrated by Berman and Conteas
(1978). These Doppler noise records offer both a means of determining the
general level of scintillation and a means of determining the index. In the
search for evidence of streams of plasma from the sun, these should serve asprimary data because they will be available for many more hours than any other
data form. Records such as Figure 8 should be available whenever the direct
recordings exist. As we will shortly illustrate via tabulation, these are
relatively plentiful.
*
Now entitled "The Telecommunications and Data Acquisition Progress
Report."
194
2 >
ccc
4~o U. *04c
>00
>OOg d i
zU
>01cc cc 2
-j t a i. 1 11
43. ... to .
0 *- In L- I..
UQ 0
-A -
cc'
g4 u
2 2
C (L 0 o4
z- z
04 0 cc43.
I w
z 0
195
10,000
00
'E.R
.E:R
)-100
zI~lz
L 100z
DAY 247
10 200 400 600 800 1000 1200 1400
TIME - minutes
Figure 8. An example of one day's "Doppler noise" records from Pioneer 11,three days after its Saturn encounter. The curves were derived from thesame parameter, but averaged over 0.1, 1, 10 or 100 seconds,
196
3.1 I)oppler Noise vs Electron Content
Figure 9 is provided by A. L. Berman (Private Communication) as anillustration that the Doppler noise varies roughly in proportion to the elec-tron content; that is, to the average plasma density along the signal path.In this record, the electron content was derived along the path to Mars usingthe dual-frequency signals from Viking. The differential group delay approachwas used and so the total content could be determined. From the coordinatedvariation between the Doppler noise and the content, plotted together here,one can see that there is indeed a strong correlation between the two. Theyare not completely coupled, however. The Doppler noise is, in this respect,much like the scintillation index in that a high value represents a combinationof the prevalence of high speed, high density, or high turbulence in someunknown combination. From the standpoint of physics, the observation providesa measurement of the inhomogeneity of the refractive index distribution; thisinhomogeneity can be brought about either by a dense region of relativelysmall turbulence or a diffuse region of intense turbulence.
When direct recordings are available, it is often possible to determinethe differential phase delay. The differential phase delay cannot bedetermined from the Pioneer Saturn direct recordings, because it only uses asingle frequency. The Voyagers have a capability to modulate both the S- andX-band signals providing the information necessary to calculate differentialgroup delay. Putting these facts together, we see that direct recordings fromPioneer Venus provide differential phase delay capability and direct recordingsfrom Voyagers provide both differential group and differential phase delaycapability. None of these data have yet been processed. Because of intensescintillation we expect that the closed-loop receivers often have detectedsignals that are too unstable to support these differential delay calculations.Until we have had an opportunity to attempt the processing, we cannot ascertainhow much of the electron content measurement can be obtained by these directmeans.
3.2 An Example of Past Interplanetary Correlations with (ontent
To illustrate the value of electron content, and to stress the importanceof correlative analysis, Figures 10 and 11 show a series of events that wereobserved in early 19b9. The relative positions of the earth and of the threeparticipating spacecraft are sketched in Figure 10 from which it can be in-ferred that the solar wind disturbances causing the correlated periods ofactivity must have spanned many tens of miliions of kilometers. The spiralsshown in Figure 10 depict the steady-state flow pattern from the sun basedupon the Carrington rotation rate and an assumption that the solar wind speedwas 400 kilometers per second. One dot is shown for each one hour of travel.The path from earth to Pioneer 8 is shaded to emphasize that the dual-frequencysystem measured the average solar wind density along the path. (This was aStanford University-Stanford Research Institute* experiment or four Pioneers.)The density is the electron content divided by the length of the path. Asshown on the top plot in Figure 11, the electron content underwent about sixdiscrete increases, which correlated one-to-one with clearly recognizable
Now SRI International
197
Hilllil I1 I I11 I 111
00
0 0 :~.. . 0
0 &J
41,
z - - .0.**
m 0 w 4
~~0.
ccv
.j.
'A 0 41
4j14
S4.I
. . . . .. . . MO
S..: 4
IN31N03 N0&13313 WO MsONa Wfldd00
198
Path fromEarth to Pioneer 8
.Venera 6 ..
Pioneer 9
Solar Wind Spiral Pattern
Osun
Figure 10. Positions of Earth and of three interplanetary spacecraftdutring the interval depicted in the next figure.
199
- >
z M
0~~ 0 C
E S 33
Co ,
* '| l
200
... .... . . .- m n mulh li l a l
,vents at earth, at Venera b6, and at Pioneer 9.
The Soviet data, comprising the Venera flux and the earth current index,Ed, from Borok, USSR were obtained from Gringauz et al. (1970). The Pioneer 9record of broadband wave level is that of Scarf and Siscoe (1971). Scarf andWolfe (1974) have shown that these wave levels correlate with the solar winddensity in the vicinity of the spacecraft. The events shown in Figure II,interpreted in the light of the geometry of Figure 10, are intended toillustrate that the content observations may yield valuable information ifthey are correlated with other measurements. The likelihood of correlationis reasonably high, as is shown by the reliable appearance of the same sixregions of activity at roughly the same time intervals at the four widelyseparated interplanetary locations.
3.3 Linear ltecorins
Processing of linear recordings by a conventional digital computer wouldbe prohibitively expensive; fortunately, we have an "array processor," which
i a specially wired computer adapted for the Fast Fourier Transform. At thetime of this writing we have working programs that have produced sample data
for the Narrabri meeting, but results of scientific value are not expecteduntil 198U. Figure 12, kindly supplied by the Jet Propulsion Laboratory(Holmes, Private Communication, 1979), shows spectra like those that are
presently being derived by our software from the available tapes.
By programming the computer to simulate a phase-lock-loop receiver, it ispossible to carry out processes that are unattainable by means of electronics.The signal is detected and its radio frequency phase is recorded for subsequent
statistical manipulation. Most of the useful information relating toscintillation is derived from this phase, although for some purposes the
strength of the signal (which also scintillates) provides a useful data source.For the moment, we can report only that the linear recordings exist in
adequate quantity and their quality seems to be excellent. The processing isjust beginning and steady progress is anticipated.
3. I ( )I her I'arl iFipaZ . (;roiii
In addition to the study of these linear recordings by array processorConduc td by our group at SRI International, another group at nearby StanfordUniversity, under the direction of Professors V. R. Eshleman and G. L. Tyler,is engaged in a more extensiv study of the Voyager linear recordings.Additionally, Drs. R. T. Woo and J. W. Armstrong, at the Jet PropulsionLaboratory, are studying the linear recordings from Pioneer Venus and PioneerSaturn. The Stanford and Jet Propulsion Laboratory groups are also utili .ngthe closed-loop records (D) as an adjunct te their studies.
201
TYPICAL SPECTRAVOYAGER 2 @ SEP = 0.6 (MINIMUM)DAY OF YEAR 227 (15 AUGUST 79)
S-BAND X-RAND'01:45uTc! 0T4 *
L__ ~ ~ W J___.
0 F 1KHz 0 F 3 KM:
-T
.. . .... .. 7 T - , z -
0 00 F I KHt 0 F 3 KHz
oS:331 UTC 05UoC'
0 F 1 Kz 0 F 3 KM
Figure 12. An intermediate data product: Spectra at S-band (2295 MHz)and at X-band (8415 MHz) derived from linear recordings of the Voyager 2signals (D. P. Holmes, private communication, 1979).
202
4. SCHEDULE OF DATA AVAILABILITY
Most of the information about the Pioneer data is now available, butthe Voyager picture is still incomplete. In any case, the Voyager recordswill contain more of the closed-loop data (D) than is shown in the followingtables.
In Tables I through 8, we have included a computerized status report thatsummarizes the extent of the data. In the coming months, as the pictureemerges and more details become clear, we will make succebsive corrections tothis computer file and periodically issue a corrected set of tables to anyscientist who expresses an interest.
In addition to the part of the tables that describes the availability ofdata, we plan to add comments about the activity, evident in those data.There was (fortunately) a great deal of solar activity during STIP IntervalVII. For the moment we have not attempted to incorporate measures of solaractivity although it would be a natural extension of the record to includesuch information.
To read these tables correctly, please note that an entry of D means thatthe direct recordings are available for some part of the one hour intervalbeginning at the time shown in the heading. Similarly, L means linearrecordings exist. Typically, recordings are continuous for several hours. Acontinuous string of Ds extending over five consecutive hourly periods mightmean that the record lasted as much as five or as little as three hours: thatit began in the first hour and extended into the fifth hour. Those investiga-tors who find the first and last hours of timing critical should contact usfor greater detail. The inner hours in any such group of letters can be con-sidered as fully covered.
5. SUMMNARY
During the passage of the four spacecraft through conjunction, weobtained extensive records of radio signals that travelled the solar coronaabout 1 to 2 N of the sun. Four types ot derived data are starting toappear as the processing gets underway, but few experimentally significantresults are ready to be presented. With the exception of the Voyager closed-loop (D) recordings, we provide a summary of the schedule of data that willbecome available. Those who have relevant observations of the sun or of thesolar wind are encouraged to collaborate in the study of these data.
Acknowledgments
This work was sponsored by the Air Force Geophysics Laboratory and by theNational Aeronautics and Space Administration.
203
References
Berman, A. L., and Conteas, A. D.: 1978, "Radial and solar cycle variationsin the solar wind phase fluctuation spectral index as determined fromVoyager 1978 solar conjunction data", The Deep Space Network ProgressReport 42-48, p. 58.
Croft, Thomas A.: 1979, "A graphical summary of solar wind electroncontent observations by Pioneer 6, 8, and 9", J. Geophys. Re s., 84,
439.
EOS: 1978, "Solar wind study", EOS Transactions, 59, 906.
Gringauz, K. I., Troitskaya, V. A., Solomatina, E. K., and Shchepetnov,R. V.: 1970, "Variations of solar wind fluxes observed on board Venera-5and Venera-6 from January 21 to March 21 1969 and pulsations of theearth's electromagnetic field caused by these variations", in V. Mannoand D. E. Page, (eds.), Intercorrelated Satellite Observations Related
to Solar Events, D. Reidel, Dordrecht, p. 130.
Scarf, F. L., and Siscoe, G. L.: 1971, "The Pioneer 9 electric fieldexperiment: Part 2, Observations between 0.75 and 1.0 AU", Cosmic
Electrodynamics, 2, 44.
Scarf, F. L., and Wolfe, J. H.: 1974, "Pioneer 9 plasma wave and solar plasmameasurements for the August 1972 storm period", J. Geophys. Res., 79, 4179.
204
0N N 0
-4 -0AD 0
a'a
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In I a 0 0 0
E0 1 0 0N C, c I C1 ci00
-44 DC an -4 0 CN0 Os 0nI mp - 0 09 00
aO CC 0 im 0 CC O 0 oO -0 ~0Xc MC)3n N 03 in 03 aO Q~ O a0 ~ C
0 00 0
-0 0
C 0a'a
0 -o a 0 0a
-4-
0m Q 14 I 4 -0 00 4 0 in0q N 0 q 00 0
0 0 Ln 0 0o 0- O 0 C
0 In 00 am0a
C 0 0 200
a 09
4 0 59
sin -i0 0 a
E. - 0CD0
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Im I
50 %D 00 CD
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RE 0:C (NQ &IN0~' N
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r- coe
En o
- 206
M" on so Do min min a a- 0 09Q 00n 00 0nI n0 0n 0
am a ~ 0 0 aim 000 00 Q 0'-4 - n no 00 a-~ in
.- no 00 Q 0Q
a- 0 00 a.-IQ 0n 0
I0 In a 0 0an 0 no 0m- 0 0a 0 f 0 0 00
Q 0 In 0ain in 0n 0
0- M Q 0 0a ~ 0
---co 0 0 IM 0 q 0n Ln 0 in LA 0 LA
-D on no (no00
no 00 in
I 00 I 00 I ai00
-(NC 00 00)0
-c 00 moo Ono 00 0 a0 mo
C.)s 000 00o n0 00 0 00o 00kI in a0 0 M 00 0o0 CC)0 000Ir, 0 In a I 0n0 000 0 0
0 in 00 0M0 a00 00n 050
-4 -c 0 0 0 0 00I o 00 C 000 0Q 0 0 0 00 000 C~q a 50
00 00nQ c0 0000 5 0 50a
a00 0 a 0 000a a 5 500 Q0 0 0 0M 00O0 050 50i
00 00 Q 000 s0 00 0 0 00n
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--- 0T L 0 0oIM r 00 005
LA -- C14 s, 0 A .. ~ ~
kom~4 00 0 0 00 o 00 S00 00 00 00 no an S
-4) M CC ai
207
00
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16. Unusually Intense JovianDecametric Emission Observed
on 1979 March 7, 1920- 2040 UT
K. Bullough and W. GibbonsDepartment of Physics,University of Rheffield,
Sheffield, S37 RH. U.K.
Abstract
On 1979 March 7 in the period 1920-2040 UT extremely intense Jovian decametricemissions were recorded at Bradfield (53026
' N., 1035
' W.) near Sheffield. The
emission was first detected after dusk at 0920 UT, when radio and communicationinterference fell to a low level. The initial signal intensity lay in the rangelo
-17 to l0
-16 W m
-2 Hz
- 1, rising to 10
-15 W m
-2 Hz
-1 in the period 2011-2027 UT and
finally attaining a remarkable 10- 14
W m-2
Hz-1
at 2032-2040 UT. The emissionceased at 2041 UT. The Jovian SIll longitude and lo phase at the beginning and endof the period ware 2100.5/61o.2 and 2480.8/720.5 respectively; corresponding to anon-lo-related source of emission.
A detailed analysis of the scintillation spectrum and indices h"s made itpossible to identify contributions from the ionosphere and interplanetary medium andin the latter case to distinguish clearly between scattering from power-law andCambridge small-scale irregularities. Immediately following an abrupt increase In
209
signal intensity at 2011 UT the ionospheric Faraday scintillation virtually dis-appeared, leaving only the interplanetary scintillation component (iPS). tiisindicated (1) that the source signal was almost perfectly circularly polarized and(ti) that tile contribution to the scintillation from the Jupiter ionosphere wasnegligible. The scintillation data indicated a hpoint sourceo at this time withdiameter <2000 km and a possible diameter during the most intense emission of op to30,000 km.
It appears that this event was initiated by the same solar wind sector andparticles that 12-14 days previously had initiated a great magnetic storm at theEarth. ..This dependence on magnetic disturbance is similar to that of auroralkilometric radiation on Earth.
An extremely efficient mechanism for particle-to-wave energy conversion andhighly beamed emission are required. We suggest that the emission may haveoriginated in cyclotron solltons (Cole and Pokhotelov, 1980) on a northern fieldline where the local gyrofrequency was greater than but close to 9 Mlhz, and whosedirection at the source was towards the Earth. Such an event would he detectable atstellar distances.
I. INTROI)UCTION
lDecametric radio emission has been studied systematically for more than 20years; for reviews see Warwick (19h7), (ehirels (1976) and recent information fromthe Voyager programme (see .Jue 1979 issue O Scienie) . Thesc studies hiveestablished such features as: (i) the identification of I.-bursts and S-bursts;(if) the signal polarization (usually RH ellip ially polarized, especially atfrequencies >20 MHz, and with the major axis ilosc to perpendicular to the axis ofrotation of Jupiter; (iii) the effect of propagation through the interplanetarymedium and ionosphere on signal scintillation and fading; (iv) the dependence ofsignal occurrence on the lupiter SII I longitude and lo phase. In particular theinfluene of Jovian SIll longitude, Io phase, and earth Jeclination, Df.:, on thereception of decametri emission has been shown to implV a very highly beamed(_60 cone) signal.
It has been suggested that tile main source of emission, centred on SIll longi-tilde 240', is associated with energetic particle precipitation into the magneticanomaly in the northern hemisphere centred on Sill longitude 2700. The most commonpolarization (Ri) for simple cyclotron emission is consistent with a source nearthe North Pole. Cole and Pokhotelov (1980) have recently proposed a cvtlotron soli-ton mechanism for generation of tile airoral ki lomet ric radiat ion observedi in lit. Eart h'smagnetosphere. This process is extremely efficient for particle-to-wave energyconversion. The radiation is beamed, at the source, along the field line (RHcircularly polarized in the north) and has a frequency approximately twice thelocal gyrofrequency. It appears to be the only model which can account simultan-,tosle for tile signal power, polarization and lack of Faraday rotation and scintilla-tion in the fraction of the propagation path lying in the Jupiter ionosphere andmagnetosphere.
The observations described here were made during a project for third-yearastronomy students which included, besides the recording of Jovian decametricemission, the measurement of galactic noise and a study of the diurnal behaviourof tile ionosphere and its effect on radio reception. A more detailed paper isbeing prepared (Bullough ct al., 1979).
210
2. EQUIPMENT AND DATA ANALYSIS
The equipment (Figure 1) consisted of a three-element Yagi mounted at a heightof 0.534 X above the ground at Bradfield (53026
' N., 1035
' W.) near Sheffield.
YAGI
Racal J Amplifier Detector C hartf
Figure 1. A block diagram of the equipment.
The beam direction was mostly maintained at a bearing 155' east of north. Duringthe period of observation the azimuth and elevation of Jupiter increased from
135' to 1660 E. of N. and 51' to 570 respectively. The influence of the F-regionon radio interference, which allowed us to recognize the Jupiter signal at 1920 UT -
and relatively unperturbed from 1940 UT onwards - is shown in Figure 2. The strongsignal drove the receiver into a non-linear region of its response and a gain
adjustment was made at 2015 CT. The original record was digitized at 2.4-sintervals, and after careful calibration a graph of the corrected signal strength(dB) was obtained (Figure 3). We also analysed selected 2-min periods which werephotographically enlarged (10 times) before digitization at intervals of 0.24 s.The effective frequency response for the chart-recorded signal was 35 Hz (to -3 dB)and 70 Hz (to -6 dB) with corresponding response time, At r (rAf)
-l, of 0.01 s and
0.005 s respectively. The low time resolution was appropriate to a study of the
ionospheric contribution to signal fluctuation (-10 s) and also fluctuation in the
source. The high time resolution made it possible to study the IPS component. Adetailed spectral analysis was made and the scintillation indices S2 , S4 t tained(Briggs and Parkin, 1963; Rufenach, 1972, 1975).
3. BACKGROUND TO TIlE OBSERVATIONS
The observations were made during the ascending phase of sunspot cycle 21.However, the overall sunspot activity was far higher than predicted. Zurichsunspot numbers for January, February and March were 166, 138 and 137, that for
211
N900 60 30.-W E-..30 60 90 °
60
N
30 "3d O
- AerialMin .3eAra
Max
//// Area of ionosphere which could reflect 18MHzsignals to Bradfield on March 7, 20.00 U.T\\\\\ Possible Location of transmitters.
Figure 2. A map to show the area of the F-region which could havesupported 18 MHz transmission to Bradfield on 1979 March 7 at 2000UT, and the geographical location of possible transmitters,assuming a single-hop transmission.
dS .,)ITER RUIl G HEAN, 1979 MAKCMI, 920.2041UT BASIC INrEVAI.2.A4 COWS
--
A B C
Time after 1920 UT (min)
Figure 3. The Jupiter signal (1920-2041 UT), intensity (dB) vs time, digitizedat 2.4-s intervals and corrected for receiver response. A three-point runningmean has been applied. At 1920 UT: 0 dB - 0.72 x 10-20 W M-2 Hz -1 . At 2040 UT:0 dB - 2.03 x 10- 20 W m- 2 Hz- 1.
212
January being the highest recorded since August 1959. The sector of solar windwhich gave rise to tie March 7 event is tile same as that which gave rise to thegeomagnetic storm of February 21/22, in agreement with tle work of Barrow (1979).The dependence of Jupiter emission on solar wind parameters (velocity, inter-planetary field strength and orientation) may therefore be similar to that for sub-storm activity in the Earth's magnetosphere (Akasofu, 1979).
Jupiter was situated at a distance of 4.57 AU from the Earth at a solarelongation of 1330 E. Thus the Jupiter ray propagated closely parallel to theinterplanetary field in the vicinity of the Earth. The declination, DE, of theEarth, was +0.64. Assuming a mean height of 350 km we see that the ray inter-
sected the F-layer close to Slough (47035' N., 9041
' E.), as listed in Table 1.
Table 1.
Intersection of the Jupiter ray with the F-layer
Bradfield Angle of Distance
Time incidencc, Ray/geomagnetic fromAzim. Elev. Long. Lat. to field angle Bradfield
F-layer to layer
LI (km)
1920 135'.0 500.9 10.1 E. 510.7 N, 360.7 420.1 444
2000 1490.3 55'.5 00.1 E. 510.6 N. 33'.4 391.0 424
2041 1650.5 56.9 00.8 W. 51'.6 N. 310.2 400.8 413
The Slough ionosonde and the Kp indices for March 7 indicated a normalF-layer with no Spread F (Kp = 3-, 3, 1+, 0+, 0+, 1, 2-, 2) (fDF- (MHz) =11.9, 9.2 and 7.6 at 1900, 2000 and 2100 UT respectively). These f0F ,
values andthe Ionospheric Prediction Service MUF graphs (published by U.S. Department ofCommerce Publ. OT/TRER 13 Vol. 1.4) were used to construct Figure 2.
I. T11E SIGNAl. SCINTILLATION
From Fremouw and Rino (1973) and Briggs and Parkin (1963) we obtain a pre-dicted scintillation index S4
= 0.04, much lower than that observed for the low-
frequency end of the scintillation spectrum (Figure 4). It was found possible toattribute the low-frequency component to Faraday rotation as the F-layer faded anddifferential Faraday rotation as the F-region irregularities -loved through thefirst Fresnel zone. This deduction is confirmed, in remarkable fashion, by thevariation in source signal polarization. When the signal became almost perfectlycircularly polarized at 2011 UT, the ionospheric contribution to the scintillationdisappeared (Figures 4 and 5).
Following the method of Readhead (1971) we calculated the scattering power of
the IPS Irregularities as a function of distance from the Earth to Jupiter and
213
Figure 4. The power spectrum immediate-ly after the sudden increase in signal
intensity at 2011 UT.
hence, by numerical integration, SL, as a function of source angular diameter(Figure 6). From Figure 6(a) it is clear that the scattering occurs close to theEarth, and assuming a point source, Gaussian irregularities and weak scattering, weobtain a maximum value of S4 = 0.53. In Figure 5(a) S4 (total) falls dramatically,at the moment of signal increase, to 0.61. This is close to the value 0.63 whichwe obtain, by extrapolation, from the Cambridge 81.5 Mi0z observations, and somewhatabove our calculated value of 0.5!. Afterwards the index recovered steadily tounity, later falling to -0.8 during the period of most intense emission. The onlypossible explanation for the minimum value, S, = 0.61, i.s :That the signal becameailsost perfectly circularly polarized, thereby eliminating the ionospheric contri-bution and demonstrating, fairly conclusively, that Faraday rotation and different-ial rotation are responsible. We may also conclude that Faraday rotation andscintillation in the Jupiter ionosphere/magnetosphere arc both very small, indicat-ing that the source is located some distance above that height in the ionospherewhere the plasma frequency, I , equals ,, the signal frequency. An importantfeature of Figure 4 is that, ecause of the enormous signal power available, weobserve and can identify the IPS contribution from both the power-law irregularities( 0.l to 1.4 liz) and also the Cambridge small-scale irregularities (v > 1.4 Hz).The best estimate of solar wind speed in this part of the solar cycle is 400 km s
- 1
(see Gosling et al., 1976). For distance z 0.2 ALI and solar elongation 1I3' thisyields a Fresnel cut-off frequency \f of 0.23 Hz (Rufenach, 1974) close to thatobserved. For the small-scale irregularities which are possibly elongated alongtile interplanetary magnetic field (Cronyn, 1972) we obtain 1.17 , v < 2.11 Hz forscale sizes in the range 250 to 139 km, in excellent agreement with Figure 4. Fromthe slope (rather ill-defined) we obtain a power-law spectrum for the irregularities:P(k) k-1 where k is the wavenumber and a = 4.6 (somewhat higher than usual).
There are three independent features at 2011 UT which indicate a veryloi ized source:
(i) an extremely rapid rise of 11 dB in .0.01 s (.'. cAt - 3,000 kin);
(ii) a cohesive source giving rise to a circularly polarized signal;
214
09-
0-8 S4(tota)
07-
06-
04-
03-
02-
01T S( .1. 4Hz)
19.45 50 55 2000 05 101 15 20 2510 35 40 UT
A ' B ' C
Figure 5(a). The variation in the scintillationindex S4 in the period 1945-2040 UT. (i) Allfrequencies. S4, (total). (ii) Contribution fromthe Cambridge small-scale irregularities(S4 1.4 Hz).
0 20-
0.18-
016
0.14-
012-
010-
008-
(0 14-I 4Hz)004
0.0 -1 4 Hzf 14Hz
195505 200 t I' 1 251 301 3S 40 UT
Figure 5(b). The temporal variation in the S4. ratios:
M.1 S,(>1.4 11z) (ii) S4 (>1.4 Hz)
S,,(<1.4 Hz) S,(0. 14 v< 1.4 Hz)
215
.706S2r)- r e
F i.b " 0 , (a)
(b) Th. scintillatio index1 S8 vs2o0e nua
1
di ameter
Figure 6. (a) The scattering power, S(r), multipliedbv the correction factor, A(f,b,h) of Budden andUscinski (1970) as a function of the distance, z, toJupiter.(b) The scintillation index S4 VS. source angulardiameter.
(iii) the scattering from the small-scale irregularities which attained a maximumvalue at 2011. This implies a minimum source size at that time (Figure 6b).
Gotwols et al. (1978) have pointed out that the effect of a finite-diameter
source is to apply a low-pass filter to the power spectrum of the equivalent pointsource. For a source of Gaussian brightness distribution the temporal source cut-off frequency, vs, equals u/n30z, where u = v. sin c. In a preliminary estimate weobtain a source semi-angle, Ejo <0".3 arc at 2011-2013 UT and I" to 1".5 arc inperiod A and up to 5" arc in period C. Corresponding source diameters are <2000 kmat 2011-2013 UT, increasing up to possibly 30,000 km in period C (see Figure 3).
5. TIlE SOURCE ON JUPITER
In the previous sections we have ascertained that Jupiter had three relativelysustained levels of emission (periods AB,C) with a maximum sustained emission inperiod C of -2 - 10
-1" W m
-2 Hz
-1. From the scintillation data it appears that the
source became strongly circularly polarized and had linear dimensions <2000 km duringthe sudden increase at 2011 UT, and perhaps up to 30,000 km in period C. We havealso found that the solar-wind co-rotating magnetic sector that encountered JupiterIn this period was the same as that responsible for the great magnetic storm(Kp • 6+) which occurred on Earth some 12-14 days previously. This was a non-Io-related emission, and the maximum emission occurred very close to the S111 longitude(2 5 1 *), corresponding to the occurrence probability maximum for 18 MHz (seeFigure 2 of Thiemann and Smith, 1978).
The flux, S, of decametric radiation received on Earth, is given by
216
P I s = .Tf-* R" '(1)
where S (W m- 2
Hz-
) is the flux received at the Earth,
P (W) is the transmitter power,
Af (Hz) is the bandwidth of the transmission,y (sr) is the solid angle over which the radiation is beamed, and
R (m) is the distance of the transmitter.
For the maximum sustained power in period C (S 2 o 10 1 W m-
Hz-
and
distance R =4.5736 AU = 6.8335 - 1011 m) we obtain
P = 9.3393 × 10' Af y W . (2)
For a cone of small semi-angle 6
P = 7.1113 × 106 Af 62 (3)
Jupiter rotates through 10 in 1.65 min;
. for period C: u < 2.420, which yields
1, 4.18 ' 107 Af W
Now .,f, from previous work, is greater than 105 Hz and usually 1 MHz
•'. 4.2 - 101 2
- P - 4.2 × 1013 W.
The flux observed from Jupiter is normally
10-?? S 10- 2
n W m- ?
Hz- 1
(and may attain 10-" W m
-' Hz
- 1 for S-bursts). Therefore the observed flux is
-10 ' times that normally observed. A flux so high has never been observed before.
There are two main controlling factors which might explain this anomaly.
(i) Decametric emission occurrence is very sensitive to the value of DE, even
though the range (±30. ) of the latter is smaller than the displacement of
the magnetic North Pole from the geographic North Pole (Figure 7). Thephase of the Jupiter period (11.86 yr) relative to the solar cycle (11.2 yr)
is such as to yield negative DE and suppress the effect of enhanced solaractivity until this present cycle (21).
(ii) there was a correlation with solar wind parameters which initiated a great
magnetic storm on Earth.
Furthermore, conditions for Jupiter observation are, at any time, limited to a
particular area (Figure 2) o" Lhe Earth and hampered by radio interference,
particularly in the evening. Perhaps neglect of the non-lo-related emission in
tavotir of the supposedly stronger to-related emission has caused these events to be
missed.
217
0 0 w 0- 0
C)~ w- .0w to0
5.' 0'- . 0 .0
~~~ .. N r - 00 0
0 ~ Q CC0.. Oo .50 w on > Q,~
0 ~ E >0..0 w E
1. 0 cl
9 s... 0 .
c- o) 0.0 a' u r: r
Qk " 00 u a
00 mco -a > c. x..u 05
* 50~00 0 0
000.-' ~ 0 -- 50S0 .- .0 5.
cu0.0 ca It Q0i 0.0 0 .S 0.
c. 0 00.-0~~~ o-S '..)0
00 '-u= . .z 0c..-.
~0 r0 .
0 u x r cL m ,
:3D' 0 a.00 --
c a 0 .. 0 00Z 0
W- w0 0.0 . S.o
0 L 0 055 0a a.. o) Go - w
w.5 -4 0.05- >m) 0 V)0 00'.
.0 E
a) 0 0 v 00
u 00 3t, 4A~~~ E. 00.00
,A0),. 0 W Q . 04
w.u m P.~~
The cyclotron soliton model of Cole and Pokhotelov (1980) fits the experi-mental data very well.
(i) It has extremely efficient particle-to-wave energy conversion, thus
allowing the extremely high observed brightness temperatures to be pro-
duced.
(ii) It produces high polarizations; the polarization is RH and circular when
the field line at the source is directed towards us.
(iii) Lack of evidence of any significant contribution to signal scintillationor Faraday rotation in the Jovian ionosphere/magnetosphere indicates thatthe source has to be located at some height above that at which the plasmafrequency. wp, equals 18 Ml4z. Their location of the source at wB
> 9 Mlz,
,p 9 MHz satisfies this requirement and also results in the field lineat the source being tilted more favourably, by 2'-6', toward the Earth.
On the Cole and Pokhotelov (1980) model the source would be located on theevening side of Jupiter (Figure 7) like the auroral kilometric radiation on Earth.The Voyager observations will confirm or disprove this interpretation, since for themodel to be correct there would have to be a brilliant aurora at the foot of thefield line.
Possibly all features (power, polarization, emission geometry, beamwidth) ofJovian decametric emission may be explained on tle cyclotron soliton model. Themain requirement is that the three-dimensional soliton be located on a field linetilted most favourably toward the Earth. Thus the non-lo-related source positionlies approximately on a Iine parallel to the Sill Ilongitude 2510 passing through thenorth magnetic pole. Such a line cuts the L-shells virtually at right angles alonga magnetic meridian where the L-shells attain a low geographic latitude and wherethle field lines on a particul.,r L-shell are therefore most favourably inclined tothe Earth. The lo-related sources are situated at similar locations. The link withLo is now more indirect; we visualize heating of the ionospheric plasma as thesatellite feeds power down the appropriate field line (in addition, of course, toincreasing the flux of energetic particles). Subsequently the heated plasmaprobably diffuses upward along the field line and .p approaches wB, making thecyclotron soliton process possible. This corresponds, very approximately, torelabelling th shaded region "A" in Figure 7(a) as "B", and vice-versa, and re-drawing the arrows, at right angles to the lo L-shell, in the direction of Earth.
The -vclotron sol iton model receives strong confirmation from the "null"resulI the Voyager flypast where thl signal level remained constant duringflypasL. This is due to the effect of the decreasing distance being annulled by theworsening orientation of the spacecraft relative to the emitting source. Forefficient partLicle-to-wavo energy conversion the energetic electrons within anindividual soliton have to be phase-bunched; hence the beamwidth (60) is simplyrelated to the soliton cross-sct ional area (-1)" to 10f'
m") transverse to the fieldI ine•.
6. IDTE",FTION OF DISITANT Jt PITER-LIKE, PLAN'ETS
Using the Vo yager tehlnique of comparing the RH and LH polarized signals and ana,'riil gain of 1' lar)a ( 5(1(1 ml we could attain a flux sensitivity ofI W m
- hz - and detect a1 lovian planet at a distance of 32 light years.
(arlv it is pe.rte,,tly feasible to search for Jupiter-like planets among the stars.F,,r random orlentation the majority should he detectable.
219
Acknowledgments
The authors would like to thank the members of the Space Physics Group at theUniversity of Sheffield (led by Professor T.R. Kaiser) and in particularMr. A. Strong for their help in this project. The assistance of Mr. P. Hughes,Mrs. Elaine Lycett, Mr. K. Thurston, Mrs. Diane Hallam and Mr. R.D. Cotterill wasalso appreciated.
The observations were conducted as a student project ably carried out byMr. T.L.P. Baud, Mr. S.R. Darnell and Miss Frances S. Lochtie tunder the super-vision of K.B.).
We are indebted to Mr. R.W. Smith for information from the World Data Centre,Appleton Laboratory, Slough.
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Briggs, B.H., and Parkin, I.A.: 1963, "On the variation of radio star and satellitescintillation with zenith angle", J. Atmos. Terr. Phys., 25, 339.
Budden, K.G., and Uscinski, B.J.: 1970, "The scintillation of extended radiosources when the receiver has a finite bandwidth", Proc. R. Soc. Set. A, 316, 315.
Bullough, K., Gibbons, W., Baud, T.L., Darnell, S., and Lochtie, F.S.: 1979,"Jovian decametric emissions of unusual intensity" (Preprint).
Carr, T.D., and Desch, M.D.: 1976, "Recent decametric and hectometric observationsof Jupiter", in Jupiter (ed. T. Gehrels), p. 693 (University of Arizona Press,Tucson).
Cole, K.D., and Pokhotelov, O.A.: 1980, "Cyclotron solitons - Source of Earth'skilometric radiation", Plasma Phys., 22, 595.
Cronyn, W.M.: 1972, "Density fluctuations in the interplanetary plasma: Agreementbetween space-probe and radio scattering observations", Astrophys. J., 171, LIOI.
Fremouw, E.J., and Rino, C.L.: 1973, "An empirical model of average F-layerscintillation at VHF/UHF", Radio Sci., 8, 213.
Gehrels, T. (ed.): 1976, Jupiter (University of Arizona Press, Tucson).
Gosling, J.T., Asbridge, J.R., Bame, S.J., and Feldman, W.C.: 1976, "Solar windspeed variations: 1962-1974", J. Geophys. Res., 81, 5061.
220
References
Gotwols, B.L., Mitchell, D.C., Roelof, E.C., Cronyn, W.M., Shawhan, S.D., andErickson, W.C.: 1978, "Synoptic analysis of interplanetary radioscintillation spectra observed at 34 MHz", J. Geophys. Res., 83, 4200.
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Roederer, J.C., Acuna, M.H., and Ness, N.F.: 1977, "Jupiter's internal magneticfield geometry relevant to particle trapping", J. Geophys. Res., 82, 5187.
Rufenach, C.L.: 1972, "Power-law wavenumber spectrum deduced from ionosphericscintillation observations", J. Geophys. Res., 77, 4761.
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