Group Title: Groundbased and spacecraft studies of Jupiter at decameter and hectometer wavelengths /
Title: Groundbased and spacecraft studies of Jupiter at decameter and hectometer wavelengths
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Title: Groundbased and spacecraft studies of Jupiter at decameter and hectometer wavelengths
Physical Description: ix, 216 leaves : ill. ; 28 cm.
Language: English
Creator: Desch, Michael Daniel, 1947-
Publication Date: 1976
Copyright Date: 1976
Subject: Jupiter (Planet) -- Radiation   ( lcsh )
Physics and Astronomy thesis Ph. D
Dissertations, Academic -- Physics and Astronomy -- UF
Genre: bibliography   ( marcgt )
non-fiction   ( marcgt )
Thesis: Thesis--University of Florida.
Bibliography: Bibliography: leaves 210-215.
Statement of Responsibility: by Michael Daniel Desch.
General Note: Typescript.
General Note: Vita.
 Record Information
Bibliographic ID: UF00098115
Volume ID: VID00001
Source Institution: University of Florida
Holding Location: University of Florida
Rights Management: All rights reserved by the source institution and holding location.
Resource Identifier: alephbibnum - 000178726
oclc - 03127936
notis - AAU5239


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To the family, of course.


I WISH TO EXPRESS my deepest gratitude to Thomas D. Carr, my

research advisor, for his guidance and patience and particularly

for his having entrusted to me, time and again, the responsi-

bilities of a research scholar. I have benefited in many ways

as his student.

I am also grateful to Alex G. Smith, George R. Lebo,

Kwan-Yu Chen, and Charles F. Hooper Jr. for having served as

my Supervisory Committee and for offering instruction, advice,

and council at various stages of my career. I appreciate the

kindness with which Stephen T. Gottesman and Frank B. Wood also

served in this capacity.

The hospitality shown by the Goddard Space Flight Center

personnel during numerous visits to Greenbelt was very much

appreciated. In particular, I owe much to Joseph K. Alexander and

Michael L. Kaiser, whose scientific and technical expertise in

managing the RAE satellite system was indispensable. Likewise,

I am grateful for the technical assistance of members of the

Astronomy staff of the University. Woody W. Richardson and Hans

W. Schrader have helped in data handling and photography. The

design and maintenance of the various groundbased installations

was masterfully conducted by Jorge Levy, Richard S. Flagg, and

Wesley B. Greenman.

I have benefited, too, from interesting and informative

discussions with Richard S. Flagg, Daniel P. McGuire, Michel

A. Lynch, Andrew W. Seacord, Dolores S. Krausche, James R.

Kennedy, Robert A. Smith, and Robert A. Brown.

I am indebted to Marcia L. Rackley for her thoughtful

comments regarding the early manuscript and for her patience

and care in reducing much of the RAE microfilm.

I will always be grateful to Denise Frank, both for having

expanded my mental horizon and for even believing such a thing

was possible.

The manuscript in final form was edited and typed by

Roberta Solt. Her attention to detail and devotion to her

craft have been a delight.

Financial support for this research has been provided by

an NDEA Title IV fellowship, by a University of Florida

Graduate School fellowship, by a grant from the National Science

Foundation (Principal Investigator Alex G. Smith), and by a grant

from the National Aeronautics and Space Administration (Princi-

pal Investigator Thomas D. Carr). Computing support was pro-

vided by the Goddard Space Flight Center computing facility and

by the Northeast Regional Data Center of the State University

System of Florida. The generous support of each of the above is

gratefully acknowledged.





Jovian Decametric Phenomenology, 4.
Jovian Decametric Morphology, 12.
Principal Motivations for Conducting
this Study, 19.

The 26.3 MHz Array, 23. Recording
and Preparation of Array Data for
Analysis, 32. The RAE-1 Satellite,
34. Initial Selection and Reduction
of the RAE-1 Data, 41.

Phasing Accuracy, 52. Absolute Gain
Calibration, 58.

Orbital Phase Calibration, 75.
Absolute Gain Calibration, 80.

The Dependence of the Occurrence
Probability on Frequency, 91.
Rotation-Phase and lo-Phase
Modulations, 98. Two-Dimensional
Analysis, 107.


The Occurrence Probability
Spectrum, 117. The Flux Density
Spectrum of the Emission, 121.
Rotation-Phase and Io-Phase
Modulations, 128. Two-Dimensional
Analysis, 140.

Analysis of the 26.3 MHz Results,
148. Preparation of the 18 and
10 MHz Data, 152. Analysis of the
18 and 10 MHz Results, 154. The
Selection Effect and the Frequency
Dependence of lo Control, 161.
Rotation- and lo-Phase Behavior
and Current Theoretical Models,

Dynamic Spectra, 179. Source Size
and Location, 187. Frequency
Extent of the Non-Io Emission,
188. Correlations with Solar
Activity, 190.





Abstract of Dissertation Presented to the Graduate Council
of the University of Florida
in Partial Fulfillment of the Requirements for the
Requirements for the Degree of
Doctor of Philosophy



August 1976

Chairman: Thomas D. Carr
Major Department: Physics and Astronomy

We present a study of the low-frequency Jovian emission.

The investigation encompasses the frequency range between 450 kHz

and 26.3 MHz and is unique in two respects: (i) At 26.3 MHz an

array of 640 dipoles has provided information on a component of

the Jovian emission which is three orders of magnitude weaker

than that detected with conventional radiometer systems.

(ii) Deployment of the earth-orbiting Radio Astronomy Explorer 1

(RAE-1) satellite has made available a window in the radio spec-

trum which is not generally accessible from the ground. Through

refinements in the RAE-1 data-selection process, we have


succeeded in identifying Jovian activity at eight frequencies

between 450 and 6550 kHz.

Both the groundbased (array) and spacecraft (RAE-1,

V-antenna) systems are described, as are the procedures which

were developed to obtain their absolute effective apertures and

pointing directions.

Analysis of the data has led to three principal findings:

(1) At 26.3 MHz a distinct intensity threshold exists below

which statistical dependence of the activity on lo phase (y1 )

disappears; by comparison, the System III (X ) behavior of

the emission is well defined at all flux levels. (2) lo is

seen to exert considerable control over the spacecraft-

detected emission at frequencies as low as 2200 kHz. Modulation

of the emission in the XAII coordinate is insignificant at these

frequencies, however, unless the data are intensity weighted.

The \ morphology is thus seen to evolve markedly between

26 MHz and 6550 kHz, whereas Io control is invariant with

respect to frequency. (3) Jupiter's power spectrum is found to

peak at about 8 MHz with the low-frequency shoulder less sharply

inclined than the high-frequency end. The shape of the occur-

rence probability spectrum of the emission is quite similar to

that of the power spectrum.

The Io morphology documented in (1) and (2) is seen at

once to challenge the widely held belief, based upon the analysis

of data recorded with conventional radiometer systems, that

Io control is a pronounced function of frequency. Specifically,

(1) is inconsistent with the distinct tendency for the high-

frequency (u 25 MHz) emission to occur over only a very limited

range of Io phase, and (2) is inconsistent with the low-frequency

(u S 15 MHz) morphology which exhibits little dependence on yo

Re-examination of conventionally recorded 18 and 10 MHz data

has led to a uniquely compelling solution to this dilemma. We

conclude that (1) and (2) are fundamentally correct, implying

that Io control is independent of frequency, at least to first

order, from 40 MHz down to 2200 kHz. Further we show that con-

ventional recording of data leads to a specious correlation

between Io control and frequency because of a selection effect

resulting from a nonuniform intensity sampling of the emission.

The lI and y1 morphologies developed herein are examined

within the framework of recent theoretical models, several

aspects of which we refine in accordance with post-Pioneer

concepts of the Jovian magnetosphere.

We conclude by exploring four specific areas of future

interest in which theory and experiment may complement and

reinforce one another. Relevant and definitive investigations

are proposed involving the high-resolution dynamic spectra,

Pioneer 10 and 11 data, and several of the experiments to be

included onboard the MJS-77 spacecraft.


The solar system consists of the sun,
the planet Jupiter, and some debris.
--Isaac Asimov

IT HAS BEEN CALLED A RENAISSANCE. As grandiloquent as this state-

ment appears, it is nonetheless undeniable that the astronomical

community has been virtually transfigured by a large-scale re-

surgence in planetary studies. The concept of the solar system

as a static and inert milieu has been supplanted by one in which

the sun and its ever-attendant convoy interact dynamically on

many scales in space and time. The solar system is now often

thought of as both a microscale astrophysical and macroscale

biological laboratory, a bubbling cauldron of wave-particle and

molecule-molecule interplay. The astrophysical manifestations

of these interactions are quite numerous, sometimes startling,

and often inexplicable. On the other hand, exobiology has re-

mained largely conjectural, but is no less exciting in its impli-

cations for mankind. Definitive in vitro experiments on Mars and

Jupiter will radically alter its tentative status.

It is of interest to speculate as to the source of this re-

awakening. Certainly the original suggestions by Urey and sub-

sequent experiments by Miller (1953) provided the foundation for

the theoretical analysis of amino acid synthesis in reducing

(proto-earth and Jovian) atmospheres. The ambitious Mars-Viking

Project owes much to the groundwork established by Urey and

Miller. In the physical domain, however, the existence of a

single landmark occurrence, similar to that of Urey and Miller's

experiments, is less evident. It may in fact be that a single

cause is lacking; the present scenario has evolved somewhat

slowly. Nevertheless, the development has no doubt been stimu-

lated by various factors, among them the discovery of the Van

Allen radiation belts (Van Allen et al. 1958), the first dis-

cussion of an open (i.e. interactive) magnetosphere (Dungey

1961), the verification of the solar wind by Mariner 2 (Neuge-

bauer and Snyder 1962), and of course, the dawn of the space

age itself (1957). The reader can no doubt supply several of

his own historically significant catalysts depending upon per-

sonal bias.

A startling astronomical discovery occurred in 1955, however,

which preceded all of the above events. Perhaps occupying that

unique historical niche is the serendipitous discovery by Burke

and Franklin (1955) of the intense radio emission from Jupiter.

In contrast to the solar radio emission (Hey 1946) which had

been anticipated as early as 1894 (Lodge 1894), the existence of

nonthermal planetary radiation was entirely unexpected. The

accidental, and even belated, nature of the discovery is not at

all surprising in view of the conceptual framework within which

solar system astronomy was being applied at the time. That

framework was shattered in 1955, and the distinction between

star and planet had begun to blur.

Since the initial report by Burke and Franklin, the discovery

and subsequent explication of many planetary and geophysical

phenomena has taken place; but ironically, a completely success-

ful understanding of the Jovian decametric phenomenon is lack-

ing. In the present work various experimental studies have

been conducted, both from the ground and via satellite. It is

hoped that the convergence of the various emission models toward

a single focus will be assisted by the interpretations pre-

sented here. Toward this end and in anticipation of the pre-

cise motivation of the present work, we present a brief summary

of the rich phenomenology and morphology of the Jovian deca-

metric radiation which has accumulated to date. Much more ex-

tensive treatments may be found in reviews published by Douglas

(1964), Warwick (1967), and by Carr and Gulkis (1969). Carr and

Desch (1976) and Smith (1976) have critically reviewed the

current posture of the observations and emission models re-

spectively; they have further complemented their discussions

with the relevant Pioneer 10 and 11 results.


Jovian "noise storms," as they are called, often last anywhere

from a few minutes to several hours, with the storm intensity

roughly proportional to the storm duration. Individual storms

are separated by periods of time which are entirely devoid of

activity; these interludes may last for only a few hours or they

may persist, during unusually quiet conditions, for as long as

several weeks. Noise storms will often commence and conclude

rather abruptly, although the main storm is sometimes preceded

by a brief, relatively feeble precursor.The storm intensity

itself is always quite irregular, displaying radical amplitude

fluctuations in times as short as a few milliseconds or as long

as half a minute. The individual intensity fluctuations within

a storm are termed bursts and, when ordered according to their

durations, are generally observed to be of two distinct classes.

The L bursts, which are decidedly the most common, usually vary

with periods longer than a few tenths of a second. When moni-

tored aurally, L bursts resemble the sound of waves breaking

upon a beach. The S bursts are characterized by much shorter

durations, typically between 1 and 50 msec each, creating an

With the proper equipment, a single noise storm can reveal
nearly all of the phenomenology to be discussed in this section.
The "morphology" of the radiation, reviewed in Section 1-2, re-
quires the long-term accumulation of data and is thus of a
more statistical nature.

unforgettable staccatolike chorus. Repeated transitions between

L and S bursts within a single noise storm are not uncommon, and

at such times the two types of bursts are sometimes recorded


The gross phenomenology of Jovian noise storms described by

the foregoing can be, and indeed has been, gathered by the most

modest of radio astronomical procedures. It is thus subject to

a plethora of qualifications and refinements to be briefly

sketched below. Nevertheless, it is worthwhile bearing in mind

that the sporadic nature of Jupiter's noise storms and their

inherent unpredictability are among the most characteristic

features of the low-frequency radiation.


The phenomenology revealed by the dynamic-spectral studies is

highly detailed. The early investigations at low resolution

first uncovered the frequency drifting of the noise storm en-

velopes (Warwick 1963). Drift rates, either positive or nega-

tive, are usually less than about 1 MHz/min. The sign and to

a certain extent the magnitude of the drift are determined

predominantly by the Jovian longitude facing the earth at the

time of emission (Dulk 1965). Noise storms have been observed to

attain a frequency of 39.5 MHz, but then only rarely. As ob-

served from the ground, activity is not usually recorded below

about 10 MHz because of the absorbing effect of the earth's

ionosphere. At resolutions exceeding about 0.1 sec and 50 kHz,

the structure of the L and S bursts themselves becomes apparent.

The L burst emission envelopes, as they appear in the frequency-

time plane, are often streaked by tilted, parallel bars which

are nearly devoid of activity (Riihimaa 1970). These "modulation

lanes" have varying slopes which are dependent upon Jovian

longitude, indicating that they are either intrinsic to the

emission process or are imposed on the envelopes during propa-

gation through the planet's magnetosphere. On the other hand,

as revealed by spaced-receiver studies, the envelopes them-

selves are apparently generated by drifting inhomogeneities in

the solar wind plasma (Douglas and Smith 1967). The implication

here is that at least a certain class of the L bursts is in-

trinsically long enduring, perhaps as long as several minutes,

but with the characteristic 1- to 30-sec modulation imposed by

interplanetary "obstructions." At still higher resolutions in

frequency and time, the S bursts (Figure I-1) begin to display

a wealth of structure and potentially a great deal of informa-

tion. While they occur rather infrequently, the S bursts are of

particular interest because they are believed to be manifesta-

tions of one of nature's more bizarre phenomena--the Io effect--

to be discussed below. A possible hint as to the nature of the

mechanism which couples Io to the source region near Jupiter

follows from the observation (Ellis 1974) that keV electrons

of a specific pitch-angle distribution may be responsible for




26 55

26 05

26 55

26 05

(a) 19 APRIL 1971 0939 0940 UT

(C) 23 MAY 1972 0839 0840 UT 50msec -
"" t-50msec-l

FIGURE I-1. High-Resolution Dynamic Spectra of Jovian S bursts (from Krausche et al. 1976).

(b) 9 MARCH 1973 1548- 1549 UT

the generation of the simplest of the S-burst classes (Figure

I-la). Proper interpretation of the more elaborate features dis-

played by some S bursts (Figure I-lb,c) is a more formidable

problem which is just now being undertaken (Flagg et al. 1976).

It is becoming clear now, however, that the separation of bursts

into L and S classes, while appealing in its simplicity, is cer-

tainly inadequate.


The average power spectrum of the Jovian radio emission is

illustrated in Figure 1-2. In order to properly compare the

sporadic decametric component with the continuous decimetric

and thermal components as shown here, Carr et al. (1964) com-

puted probability distribution functions for the peak flux den-

sities occurring in consecutive 10-min intervals. In this way

a mean power at each frequency in the interval between 10 and

27 MHz was attained. It is of considerable importance to note

the extremely steep nature of the decametric portion of the

spectrum. Radio emission occurring at 27.6 MHz is over two

orders of magnitude less intense than that at 10 MHz. Moreover,

if the peak (rather than the mean) flux densities at each fre-

quency are compared, the behavior is quite similar; the spectrum

is merely shifted upwards by three orders of magnitude. As will

be discussed further in Section 1-3, this aspect of the phenomen-

ology is capable of imposing a selection effect upon the data

which could improperly bias much of our understanding.


105 I I I

-- )4 decametric

Ld 103- A

a *
X 102 thermal
LL decimetric
10 -
.... . ... I , . ... . . .... m l ,
10 100 1000 10000



FIGURE 1-2. Average Power Spectrum of the Radio Emission from
Jupiter Showing the Decametric, Decimetric, and
Thermal Components. Flux density is expressed in
Jy (10-26Wm-2Hz-1). Square points are from Carr
et al. (1964), triangles from McCulloch and Ellis
(1966), and circles from various sources, as
quoted by Roberts (1965) and Warwick (1970).


The precise location of the Jovian decameter source remains

obscure to this day. That it is within the corotating portion

of the planet's magnetosphere, that is within about 20 R stems

from the consistent repeatability of the dynamic spectral

features and from the observed stability of the rotation period.

In fact, there is evidence that the source is much closer still,

say within 1 to 2 R of the planet's center. For example, an

appeal for a source radiating at the local gyrofrequency ensues

from the near coincidence of the maximum observed frequency of

the noise storms (39.5 MHz) and the peak gyrofrequency at

Jupiter's high-latitude cloud tops as derived from the in situ

measurements of Pioneer 11. Gyrofrequency emission in the 10 to

40 MHz range, then, would all take place within a source region

not exceeding a distance of about 1.7 R from the planet's center.

While the arguments for the source location are somewhat

heuristic, limits on the size of the source itself have been

obtained by more direct procedures, namely from the aforemen-

tioned S burst dynamic spectra and from very long baseline

interferometry (Lynch 1972). Both methods have emphasized the

imponderably small linear dimensions of the source, estimated

between 50 and 300 km. This is particularly striking in view of

the fact that the decametric radiation, while quite comparable

to solar Type III activity both in character and intensity,

originates from a source region three to four orders of magnitude

smaller in size than that of its solar counterpart! The extreme

brightness temperature so imposed on this small emitting region

implies that a highly efficient, induced (as opposed to spon-

taneous), coherent plasma process is taking place. This assump-

tion has provided a conceptual basis upon which nearly all of

the more recent decametric theories have evolved.

Besides requiring a coherent process, it has been common to

assume further that the radiation is generated near the local

gyrofrequency. Observations of the polarization characteristics

of the waves have lent credence to this latter assumption.


At frequencies above about 18 MHz the radiation is predomi-

nantly right-hand elliptically polarized. Careful examination

of the magnetoionic theory of waves propagating in a magneto-

plasma reveals that eliptical polarization is characteristic of

radiation generated near the local electron gyrofrequency. How

near depends to a large extent upon the ratio of the plasma

density to the magnetic field strength at the point of emission.

Strict limitations cannot be established at present; however,

gyroemission appears to be a relatively safe working assumption

in view of the near coincidence of the maximum observed fre-

quency and the electron gyrofrequency corresponding to the

field strength measurements. The predominance of the right-hand

sense implies further that the radiation propagates in a

direction for which the field has a parallel rather than an

antiparallel component, assuming the extraordinary base mode.

The implication is that the source of the right-hand polarized

radiation lies in Jupiter's northern hemisphere.

At frequencies below about 18 MHz the radiation can be

left-hand or right-hand polarized, depending upon central

meridian longitude, and can be circularly polarized (Kennedy

1969). Whether this is a propagation or emission dependent

phenomenon has not been ascertained.



The long-term accumulation of Jupiter data has revealed

that there are distinct occurrence probability increases during

times when particular Jovian longitudes are facing the earth.

It was observed (Douglas 1961) that the occurrence probability

peaks were maximized when the activity was plotted as a func-

tion of the System III longitude (A1 1957.0), corresponding

to a rotation period of 9 55m 29137; presumably this is approxi-

mately the rotation period of the magnetic field and hence,

according to more recent theories,! of the metallic hydrogen core

of the planet. The discovery by Bigg (1964) that these specific

longitude regions or "sources" are further enhanced, both with

respect to the probability of their occurrence and overall

intensity, when Jupiter's innermost Galilean Satellite, lo, is

in particular locations of its orbit, was one of the more impor-

tant astrophysical discoveries of the past several decades.

Order of magnitude increases in the occurrence probability

are noted when yvl (departure of lo from superior geocentric

conjunction) is equal to approximately 900 or 2400. A repre-

sentative plot showing the locations and designations of these

sources in the 11,-vo plane appears in Figure 1-3. In addition

to being classified as lo-related or non-Io-related, each

source possesses distinguishing characteristics which are

summarized in Table I-1. It is notable that the S burst

activity discussed earlier is generally only associated with

the Io-related sources.


Following a decade of continuous monitoring of the planet, it

became evident that the position of Source A was drifting with

respect to the System III longitude defined above. With the

adoption of a revised rotation period (Gulkis and Carr 1966),

this systematic drift of Source A toward higher longitudes

was converted into a periodic oscillation of the source cen-

troid about a mean value. The long-term drift had been removed.

The period of the oscillation appeared to be more closely

aligned with Jupiter's orbital period of 11.9 yr rather than

with the approximate 11-yr semiperiod of the solar cycle. This

orbital-period oscillation also manifests itself as a periodic

90 o-a .

Io-D -

180 non'-o-B ...


270 -

I I .-" I I -

90 1800 2700

FIGURE 1-3. Representative Plot in the A --Yo Plane of the
Generally Recognized Jovian Sources. Each line
traces the instantaneous y1 -I phase during a
single storm. Heavy lines indicate activity
exceeding 6 x 105 Jy.

TABLE I-i. Some Characteristics of the Jovian Source Regions at the Upper Frequencies.

XI (1965) lo Frequency Predominant
Spana Spana Range Polarization
Designation (Degrees) (Degrees) (MHz) Sense Other Characteristics

Io-A 195-285 220-260 14-36 RH Often obscured by non-Io-A during
years of positive DE L bursts; S

195-285 0-360 11-28

95-195 65-110 11-39.5




at least Prob-
16-26 ably



0-360 at least

(at least
23 MHz)

Highest occurrence probability of
all sources, for frequencies be-
tween 15 and 25 MHz and during
years of positive DE. Greatest
change in source position, width,
and occurrence probability with
change in DE.L bursts only.
Highest intensity bursts (both L
and S). Principal S-burst source.
Most predictable source. Bifurca-
tion apparent in high-resolution
Very weak, but relatively high
occurrence probability. L bursts;
S uncertain.
Predictable. Other major S-burst
source (also L bursts).

Moderately distinct from non-Io-A.
Abundance of LH polarized emission
relative to non-Io-A may be dis-
tinguishing characteristic.
Probably L bursts only.






TABLE 1-1, continued:

AII (1965) yIo Frequency Predominant
Spana Spana Range Polarization
Designation (Degrees) (Degrees) (MHz) Sense Other Characteristics

LH Very low occurrence probability.
Characteristic dynamic spectrum
and LH polarization sense
important in identification.

aBetween approximately 15 and 25 MHz.





variation in the overall occurrence probability of Source A.

Both effects are in phase and are consistent with a model

(Carr 1972) which depends upon the changing viewing geometry in

the ecliptic plane between the earth and Jupiter. The parameter

describing this geometry is called DE (Jovicentric declination

of the earth), and the maximum excursion in either direction

north or south is only 3 5. As remarkable as it seems that such

a small angular deviation can have readily noticeable effects,

it is perhaps to be expected in view of the narrow-beaming

properties of coherently radiated emission.

Additional DE-dependent phenomena have been found, notably

the variations in source centroid positions with respect to the

yo coordinate (Conseil 1972; Lecacheux 1974; Thieman et al.

1975). It now appears that the lo-related source centroids vary

both in the A and y coordinates. One might anticipate that

the XIA and Y, variations would be related so as to effectively

minimize any changes in the sub-Io longitude (X ) on Jupiter

which presumably stimulates the emission. This is not so. Sur-

prisingly, source positions with respect to Ao change radically,

increasing by as much as 150 (for the Io-B source) as DE

changes from 00 to 3.50. Even purely geometrical explanations

of such viewing-angle phenomena are complicated by the magnetic

field models of Jupiter's intricate surface field, now emerging

from analyses of the Pioneer 11 data (Acuna and Ness 1975; Smith

et al. 1975). Observations of the emission made by spacecraft

well out of the ecliptic plane will surely be revealing in

this regard.


As implied by the above, if the apparition-by-apparition his-

tograms are plotted in the old (Epoch 1957.0) System III, the

observed drift of Source A may be used to calculate the rotation-

rate error effecting the drift. This is applied as a correction

to the .old period to obtain a new rotation period which, as men-

tioned above, insures the absence of long-term drifts. In cal-

culating the error, histograms are chosen which are separated

by approximately 12 yr to remove any bias in the source cen-

troids due to DE modulations.

The Jovian rotation period may also be determined by using

sophisticated power spectral techniques (Kaiser and Alexander

1972). Both methods, though, have resulted in similar values,

the average of these and several other methods yielding an

agreed-upon figure for the rotation rate of 870 536 per

ephemeris day. This figure, designated System III (Epoch

1965), will soon be adopted officially by the IAU. The

corresponding rotation period is 9h 55m 29s71, which should be

used when analyzing data spanning more than a few years time if

consistent results are desired. In the present work the old

System III (1957.0) will be used,as detailed features in the

1 coordinate will not be examined and the data do not span

long periods of time. Where confusion might arise, however,

differences in source locations between the two systems will be

specified. For example, comparisons between source locations as

appearing in Table I-1 (Epoch 1965) and sources discussed in

the present work (Epoch 1957.0) may be made through the simple

application of the formula

I111 (1965) = A11 (1957.0) 0.008284 At (1-1)

where At = T 2438761.5 and T = epoch of date (JD).



Nearly all of the observations of Jupiter's decametric emis-

sion have been made with low-gain Yagi antenna systems. The ver-

sion of the Yagi monitoring system which consists of a two-

element interferometer is probably the most sensitive Yagi

radiometer currently being employed in the synoptic monitoring

of the planet. Typical sensitivities (minimum flux density

necessary for the positive identification of Jovian activity)
have been quoted in the literature. Estimates range from 10 to

105 Jy (Jy = 10-26Wm-2 Hz' ) at frequencies in the 10 to 30 bMlz

range (Bozyan et al. 1972; Alexander et al. 1975a).Notwith-

standing the fact that much of the Jovian emission which occurs

is orders of magnitude more intense than this, one must consider

what morphological changes in the emission might become apparent

under conditions of greatly increased sensitivity. Such high-

gain monitoring should be carried out at a relatively high

decametric frequency for a number of reasons. The large antenna

arrays necessary for highly sensitive monitoring become pro-

hibitively large and expensive as the frequency decreases;

thus, the first reason for choosing a high decametric frequency

is purely a pragmatic one. A more important justification, how-

ever, stems from the observation that the Jovian decametric

flux spectrum falls off extremely sharply at frequencies above

about 10 MHz (see Figure 1-2). A uniform sampling of the radia-

tion as a function of frequency would therefore demand con-

formity of the flux threshold of a given radiometer system

with the flux spectrum of the emission itself. Hence, a

statistically uniform sampling is assured only when, for example,

monitoring at high frequencies is conducted with extremely

high gain systems. This has not been the case. Synoptic moni-

toring has been carried out with antennas of relatively uni-

form sensitivity in comparison with the 2.5 order of magnitude

decrease in the peak intensity of the radiation which pre-

vails between 10 and 27 MHz.

Clearly, the elimination of the selection effects embodied

in the synoptic monitoring heretofore is of vital importance.

For this reason a program was established for the purpose of

continuously monitoring the planet at 26.3 MHz using a highly

directive array. The phase and gain calibration of the radiometer

system and the subsequent analysis of the data are described in

this dissertation.


Because of the opacity of the earth's ionosphere, it becomes

increasingly difficult to do radio astronomy as observations

are extended to frequencies below about 10 MHz. This situation

improved significantly after the launching into earth orbit of

the Radio Astronomy Explorer (RAE-1) satellite in 1968.

Operating in the frequency range between 0.5 and 9.2 MHz, the

minimum frequency at which effective monitoring could take place

had been reduced by at least several octaves. However, while

the galactic background radiation and bursts of solar origin

were studied intensively (Alexander and Novaco 1974; Fainberg

and Stone 1974), the identification of the low-frequency Jovian

emission remained elusive.

In order to establish the existence, or lack thereof, of

Jovian emission over the RAE-1 frequency domain, we initiated a

program designed, at the very least, to fix upper limits on

the flux densities as a function of RAE-1 frequency. The study

was anticipated to be of considerable value as many of the RAE-1

receiver channels covered frequencies at which Jupiter either

had never been detected or had only been tentatively identified

from the ground. The requisite elements of such a study

necessarily include (i) the optimum presentation of the

satellite data for inspection and analysis, (ii) the isolation

of periods of time relatively free from known sources of inter-

ference, (iii) the determination of radiometer sensitivity as

a function of frequency, and (iv) either the establishment of

the aforementioned upper limits or the presentation of the mor-

phological features of the designated Jupiter activity. These

elements are fully discussed in the following pages.

Although the experiments described here are being conducted

at virtually opposite ends of Jupiter's low-frequency spectrum,

we will examine the spacecraft-derived results in the framework

of the morphology derived at 26.3 MHz to determine whether a

self-consistent explanation of the various phenomena is possible.


One reliable observation is worth
a thousand models and a million
speculations. --H. E. Landsberg

II-1. THE 26.3 MHz ARRAY

Located at the University of Florida Radio Observatory (UFRO)

is one of the largest, filled aperture, low-frequency antenna

arrays in the world. Approximately one half of the 640 half-

wave dipoles comprising the array are visible in Figure II-1.

It operates at a frequency of 26.3 MHz (A = 11.4 m) and has an

effective bandwidth of about 0.5 MHz. The individual dipoles are

oriented east-west (ew), and they fill a rectangle, the major

axis of which is oriented north-south (ns). The rectangular

collecting area of the 26.3 MHz array (referred to hereafter

as simply the "array") has an ew by ns extent of 16 by 40

dipoles. Since the dipoles are separated from one another by

0.6 A in both the ew and ns directions, this corresponds to a

perimeter of 9.0 A by 23.4 A. The principal lobe half-power

beamwidth (HPBW) of the antenna is approximately 600 ew by

205 ns.

A View Facing North of the 640-Element Dipole Array Located at
the University of Florida Radio Observatory. The small white
box is being opened by graduate assistants in order to make a
coarse north-south (CNS) phase adjustment. The large white
houses contain the Butler matrix networks.



I r; .L.5

I ~ -AI .

''p I

Y14 :Ej

*-f ,R' -

i- I; .


~; 1 ,L ..t N~ ih~

* 1. ~ *~ p "'

Tracking of sources across the sky is achieved by means of

Butler matrix phasing (Butler 1966). The Butler matrix is a

fairly complex network of hybrid rings which is capable of

synthesizing a number of independent antenna beams equal (in the

most efficient case) to the number of independent inputs. The

16-port (8 input, 8 output) Butler matrices employed in the

array system (see Figure II-2) generate eight independent beams

in the ew plane, each beam separated by 60. The beams intersect

at the -4 db points. They are stepped through sequentially from

east to west at approximately 26-min intervals as the source

drifts across the sky. This permits about 3.4 hours of monitor-

ing of a given source at one time.

A group of eight (2 ew by 4 ns) dipoles (see Figure II-3)

is organized into a single subunit of phasing. Eight of these

subunits, organized in the ew direction, supply the required

eight inputs to the ew Butler matrix. The synthesis of these

eight subunits by the Butler matrix results in a single main

phasing unit comprised of 64 (16 ew by 4 ns) dipoles. A total

of 10 ew Butler matrices (aligned ns) are required to control

the array's complement of 640 dipoles.

A limited amount of adjustment of the position of the main

beam in the ns plane is necessary during each step to a new ew

Butler matrix output. This is accomplished by two (5 input,

8 output) ns Butler matrices, one for each ns half of the array.

They are identical in design to the ew Butler matrices except




Schematic of the Butler Matrix Circuitry.
Similarly lettered junctions on the hybrid rings
are connected directly together. The hybrid rings
are constructed of 75 ohm, RG59 coaxial cable.
The cables connecting the rings are 50 ohm, RG58.
Each network, of which there are 10 for ew phas-
ing and 2 for ns phasing, is enclosed in its own
housing or "Butler matrix box." The control re-
lays associated with the networks are also en-
closed. In Figure II-1 the Butler matrix boxes
are the dog-house sized structures running ns
along the center of the array.

FIGURE 11-2.


hybrid balun

matching twinlead
twinlead \50 network

x.... CNS phasing cable

FIGURE 111-3.

Schematic of a Group of 8 Dipoles Constituting
the Basic Subunit of Phasing of the Array. The
dipoles, numbered D1 through D8, are not shown.
The matching of two balanced, 78-ohm dipoles
into the unbalanced, 50-ohm (RG58) coaxial
cables (which connect the hybrid rings together)
is accomplished by means of a matching network-
balun-hybrid ring assembly (see detail in figure).
Each assembly is enclosed in a small box, three
of which are color coded for CNS phasing identi-
fication. The boxes are connected together by
means of underground cable indicated in the
figure by light wiggly lines.

that three of the inputs are terminated in 50-ohm loads. Three

of the eight outputs of each ns Butler matrix are monitored at

any given time. They are separated by about 2.6 in the ns plane.

The central beam is maintained on the source, and the two

adjacent beams are used for the purpose of discriminating

against interference. The signal from the entire array is

finally arrived at by combining the outputs from the two ns

Butler matrices through a hybrid ring unit containing a variable

phase delay. There is one such unit for each of the three beams.

The foregoing phasing hierarchy is illustrated schematically

in Figure II-4.

While hour angle coverage is limited to within +1.75 hours

of meridian transit, sky coverage is complete in terms of the

declinations observable at +300 latitude (-600 to +900). How-

ever, because the ns HPBW of a 64-dipole phasing unit is only

about 240, sky coverage in the ns plane would be restricted to

this narrow range if ns phasing between individual dipoles were

not implemented. To obtain total sky coverage, then, the beam

is phased in steps of about 120 by means of manually inserted,

plug-in phasing cables (see Figure 11-2). This is referred to

as coarse north-south (CNS) phasing and is required whenever a

change in position of the beam by more than about 200 in the

ns plane is desired.

Combining the output of five ew Butler matrices at a ns

Butler matrix results in a beam which is narrow enough in the

- 8-dipole configuration,
repeated seven times to the west



variable phase delay
(SNS phasing )
* to receiver

The Array Phasing Hierarchy is Shown. Each ew
Butler matrix (designated ew) is fed by the out-
puts from eight groups of 8 dipoles each. The
two ns Butler matrices (designated ns) are each
fed from the outputs of five groups of 64 dipoles
each. The final signal is arrived at by combining
the output from each ns Butler matrix through a
single hybrid ring, whereupon the signal is
amplified in the field and fed into the observa-
tory receivers through aluminum coaxial cables
(see Appendix A).

FIGURE 11-4.

ns plane (50 HPBW) that further phasing is required for each ew

Butler matrix step, as mentioned previously. This fine north-

south (FNS) phasing is performed automatically by the ns Butler

matrices in 30 increments. The final ns phasing adjustment (SNS)

is made by the hybrid ring unit containing the variable phase

delay mentioned above. SNS values must be converted to a

rotary-switch-position (RSP) integer before being used. SNS

increments are 0 6 each. The proper delay lengths to be used

for CNS, FNS, and SNS (RSP) phasing, which are each a function

of source declination and ew beam position, are preselected

by the observer, based on numbers tabulated in the

"Phase Tracking Parameters" (PTP) booklet (Desch 1972), which

is maintained at the observatory.

The signals from the three ns beams are amplified in the

field by separate preamplifiers before being sent through 350

meters of shielded aluminum cable to the observatory. The noise

temperature of the preamps is approximately 600 K. Because of

the severe amount of loss in the shielded coaxial cables, com-

prising both the Butler matrices and all of the phasing and

transmission paths, the galactic background temperature, as

measured at the input to the field preamps, is comparable to

the preamp noise temperature itself. Considerable enhancement

in the signal to noise ratio is anticipated through the in-

stallation of ten additional preamps in order to increase the

signal from the ew Butler matrix outputs. The present sensitivity

of the array is about 50 Jy for sources near the zenith. A

factor of two improvement is expected.

In the observatory the signals from the three ns beams are

total power, diode detected by identical 250 kHz bandwidth re-

ceivers. The integrated outputs are then recorded on paper or

magnetic tape. Specifically, the Jupiter data taken in this

analysis have been simultaneously recorded on both high and low

dynamic range chart recorders. Individual events ranging from

about 0.1 db (50 Jy) to nearly 27 db (>10 Jy) above the

galactic background have been studied.


Conditions permitting, the array is operated and observations

are made beginning about 2 hours before and extending for about

2 hours after each Jovian meridian transit. Because it has been

possible to extend the monitoring program well into the daylight

hours (during that portion of the apparition when Jupiter is

near conjunction), a large fraction of each 13-month apparition

is observable.

The receiver outputs are recorded on paper, using a post-

detection time constant of 1 sec. The chart recorders are

operated at speeds sufficient to specify listening and activity

times to within +30 sec. Calibration steps are manually placed

on the recordings before and after each observing session (see

Appendix A). An observer is always present to aid in the aural

identification of suspected activity; however, the segregation

of Jupiter activity from interference is ultimately made on

the basis of the relative deflections in the three channels

described in Section II-1. For example, because station and

lightning interference generally propagate toward the ground

from a large (larger than 100 by 100) area of the sky, they

will appear with approximately equal intensity in all three

channels. On the other hand, Jupiter activity appears only in

the main channel, that is the channel which monitors the source-

centered antenna beam. Because much of the weaker activity is

completely inaudible, the three-channel identifications are

not only helpful but in fact indispensable. The accuracy with

which the array's system of beams is pointed is discussed in

Section III-1.

All data were reduced by the author, by noting the beginning

and ending times of the observing and activity intervals.

Periods of excessive interference were not counted as acceptable

listening times. The actual beginning and ending times of

Jupiter activity were recorded, as opposed to the conventional

Florida method (Register 1968) of noting the presence or

absence of activity in consecutive 5-min blocks. The single

most intense segment of each uninterrupted activity interval

was noted and scaled as to its peak and galactic-background

calibration levels. These values were converted into flux

density and normalized to a standard distance of 4 AU by means

of the gain-calibration results presented in Section III-2.

Finally, a character figure was assigned to each interval of

activity depending upon whether it displayed primarily S-,

L-, or scintillation- (single bursts enduring for more than

30 sec) type emission. All of the foregoing information was

stored on IBM cards. Approximately 500 cards (designated the raw

data deck) were necessary to compile the above information for

a single apparition. By means of Jupiter ephemeris data, the

raw data were converted via the IBM 370 computer into an

"intermediate" data deck (see Table II-). The latter contains

the radio longitudes (AII,) and satellite positions for the

beginning and ending times of the observing and activity

periods. Also stored on the intermediate deck cards were the

peak flux density and character figure for each storm. Resolu-

tion in both the I and satellite coordinates is 1.


The launching of the Radio Astronomy Explorer 1 satellite

(Explorer 38) on 4 July 1968 marked the culmination of nearly a

decade of investigation into the feasibility of making low-

frequency radio astronomical observations from above the

terrestrial ionosphere. The sounding rocket and satellite

experiments which preceded the RAE contributed immensely to the

development of a radiometer system which could function satis-

factorily in an ionized medium and be mechanically stable in

TABLE II-1. Array Intermediate Deck Card.

IBM Card Space Da ta

01-04 Station code 'F26A'
05-10 Year, month, day in format YYMDD
11-15 Julian date (Florida convention)
16-18 Begin listening Xill, 10 resolution
19-21 End listening ,,i 10 resolution
22-24 Begin activity X I', 10 resolution
25-27 End activity A l, 10 resolution
28-39 Same for Io (departure from superior conjunction)
40-51 Same for Europa
52-63 Same for Ganymede
64-66 Activity time in minutes
67-73 Peak flux density of activity interval in Jy
78 Character code: 1--S burst, 2--L burst,
79 Credibility code: 1--certain, 2--probable,
80 Coordinate system code: 0--geocentric,

spite of the antenna dimensions required for adequate spatial

resolution at very long wavelengths. As such, RAE-1 became the

first satellite devoted solely to radio astronomical studies.

The following is a brief description of the orbital character-

istics, radiometer system, and data managing technique in-

corporated in the final spacecraft design. The information has

been derived primarily from Alexander (1970) and from Weber

et al. (1971),who are members of the Goddard Space Flight Center

staff and who have worked with the RAE design from its incep-

tion. Much of the information applies equally well to the RAE-2

satellite which is now in lunar orbit.

RAE-1 orbits in the retrograde sense at an altitude of 6,000

km (1.94 RE from the earth's center). The orbit is nearly cir-

cular, having an eccentricity of less than 0.002 and a period

of 224 min. It is inclined at an angle of 600 to the equator

and processes at a rate of 0052 per day. The spacecraft alti-

tude was confined to a certain range of values by the available

launch vehicles. Within this range the altitude chosen was a

compromise between the desirability of a relatively high alti-

tude (low plasma density) orbit and a lower altitude orbit

necessary for adequate stability of the antenna structure. This

stability is derived from the gradient of the gravitational

force across the 460 m extent of the satellite and is suffi-

cient at 1.94 RE to maintain the major axis of the antenna

structure parallel to the local gravity vector. At the

altitude chosen, the in situ plasma frequencies range from

about 200 to 600 kHz.

The traveling-wave V-antennas, shown in Figure II-5, con-

stitute the primary antenna system on the RAE. Each boom is

229 m long with a 600 apex angle. As mentioned, gravity gradient

forces maintain the upper-V directed at the local zenith and

the lower-V at the subsatellite point. The antennas were

stored on motorized spools during launch and deployed in orbit.

Analysis of pictures obtained from television cameras located

on the satellite have verified the positioning of the booms

and their relative freedom from oscillation (<3 ). A 600-ohm

resistor in each boom, located three fourths of the distance

from the apex of the V, establishes the frequencies at which

the antenna is correctly terminated to allow traveling waves to

exist. This condition is satisfied at frequencies equal to 9.18,

6.55, 4.70, 3.93, 2.20, and 1.31 MHz, that is, when the section

of the antenna beyond the resistor is an integral number of

quarter wavelengths. At the frequencies for which the end sec-

tion is an odd multiple of quarter wavelengths (6.55, 3.93, and

1.31 MHz), a front-to-back ratio exceeding 10 db is achieved.

The calculated E- and H-plane antenna patterns for this ideal

case are shown in Figure 11-6. For the three satellite fre-

quencies which are less than 1.31 MHz (0.90, 0.70, and 0.45

MHz) the antenna is shorter than one wavelength and hence has

the isotropic reception pattern of a short dipole. The

upper V

lower V

A Drawing (Not to Scale) Showing the RAE Antenna
Configuration. The dipole and V antennas all lie
in the orbital plane of the satellite, with the
37 m dipole tangent to the satellite's instan-
taneous velocity vector.

FIGURE 11-5.


FIGURE -6 The Calculated E- and H-Plane Antenna Patterns at Three
Io 0"20-I 20" _10 0

Sa a o
-50 0o 3oo


FIGURE 11-6. The Calculated E- and H-Plane Antenna Patterns at Three
Frequencies are Shown. (After Weber et al. 1971.) The
E plane lies in the orbital plane of the satellite. The
H plane is perpendicular to the E plane.

upper-V/lower-V antenna configuration permits discrimination

between celestial and terrestrial noise since the signals from

each are recorded separately. The orbital characteristics allow

the upper-V to map the entire sky between +600 declination in

about one year.

Because of its relatively high directivity, the V-antenna is

useful for detecting low intensity phenomena; however, the

absolute gain of the system is not known to any great accuracy.

For this reason, a 37-m dipole was included on the satellite

since the radiation characteristics of the dipole antenna are

well known. This allows for greater precision in terms of

absolute flux density and brightness measurements. The dipole

and V-antennas all lie in the orbital plane of the satellite.

Ryle-Vonberg (RV) receivers are connected to all three of

the antenna systems mentioned above. The long-term stability and

insensitivity to gain and bandwidth changes inherent in the RV

radiometer was a determining factor in its selection as the

principal receiver system. The effective bandwidth of the

receiver is 40 kHz with a postdetection time constant of

0.1 sec. The preamplifier noise temperature is about three

orders of magnitude less than the galactic background tempera-

ture of 106 K. All nine frequency channels between 9.18 and

0.45 MHz are stepped through in 72 sec. Additional instrumen-

tation exists in the form of a fast time response, total power

receiver on the dipole and lower-V antennas. Designated "burst

radiometer" (BR), it is especially useful in generating dynamic

spectra of the intense Type III solar bursts.

The spacecraft system was originally designed to record all

data on magnetic tape to be played back to ground stations at a

25:1 speed-up. Because of a recorder malfunction after 2 months

in orbit, however, data had to be telemetered to ground stations

in real time. In spite of this, only 5% and 10% of the data were

lost during the first and second years of operation, respectively.

The raw data tapes are then edited by ground personnel and the

recorded voltages are transformed into antenna temperatures by

means of prelaunch calibration curves. The final tape record-

ings containing the antenna temperatures and all of the

necessary orbital information are then ready for computer pro-

cessing and analysis.


Although it was launched in July of 1968, the initial orbital

characteristics of the RAE-1 were such that observations of

Jupiter were not possible for the first 6 months of the

satellite's operation. As indicated in Figure 11-7, a single

RAE-1 orbit permits the upper-V to view only a portion of the

sky between +600 declination. However, observations over the

same declination range, but at successively increasing right

ascensions, are achieved through precession of the orbital

plane of the satellite at the rate of 2 min of right ascension



antenna pattern


to -600 DEC KEY:


15 14 13

12 II 10 9

Three Orbits of the RAE-1 Satellite on 1 Mar, 1 Apr, and 1 May 1969
Are Shown Projected onto a Portion of the Celestial Sphere. The
6550 kHz upper-V antenna pattern (to HPBW) is also illustrated as
it would appear projected onto the sky. The eastward motion of the
orbit is accomplished through the precession of the satellite's
orbital plane.





FIGURE 11-7.



(0052) per day. After one full year, all right ascensions between

+600 declination will have been sampled. Consequently, by January

of 1969 the upper-V beam axis (projected upon the sky) was

passing within about 450 of Jupiter when the orbital phase of

the satellite was 00. Beginning at this time and continuing for

about 6 months, detection of emission from the planet was

considered feasible; that is, Jupiter was within the V-antenna

pattern at all frequencies for some fraction of each orbit. By

July of 1969 the angle at closest approach between Jupiter and

the projected upper-V beam axis began to exceed 450, having

passed through 00 on 24 Mar 69. This period of time, from

1 Jan 69 to 1 Jul 69, defined the interval over which the RAE-1

data tapes were to be processed and examined for indications

of Jovian emission.

It was decided to process the magnetic tape data by reducing

it to a format which permitted visual identification of Jovian

activity. In addition, the large volume of data available

(over 1,000 orbits at each of 9 frequencies in 1969 alone)

required that the visual inspection of the records be carried

out at the University of Florida and, therefore, that the

records be easily portable. Both requirements were satisfied

In this study the satellite's orbital phase was arbitrarily
set equal to 00 when the angle between Jupiter and the projec-
tion of the beam axis of the upper-V antenna upon the sky was a
minimum. Orbital phase increases as this angle increases, and
the time between successive 00 transits is one orbital period
(224 min).

through the generation of 16-mm microfilm plots of the appro-

priate RAE-1 data.

An additional constraint on the observing interval, more

rigorous than that determined by the V-antenna power pattern,

was the angle between Jupiter and the sun when the upper-V

beam was scanning Jupiter. This parameter is important because

solar Type III bursts were found to be a major source of inter-

ference. Clearly, it was desirable that the Jupiter-sun angle

approach 1800 so that the upper-V/lower-V antenna configuration

might be utilized. Fortunately, Jovian opposition (24 Mar 69)

occurred precisely in the middle of the 6-month observing

period chosen for processing of the data tapes, as described

above. A 3-month interval of time, approximately centered on

the time of Jovian. opposition, was chosen as the most favorable

period over which the microfilm search was initially conducted.

The technique of discriminating between solar and Jovian

emissions based upon the comparison of simultaneous upper-V and

lower-V deflections was then optimized.

Similar considerations to those outlined above were applied

to the remaining RAE-1 data which were available during the 1970

and 1971 apparitions of Jupiter. The following table lists the

periods of time during 1969, 1970, and 1971 over which the data

tapes were processed. All of the frequencies between 450 and

9180 kHz were transferred onto microfilm. Also tabulated are

the time intervals and frequencies at which the resultant

microfilm has been read.

TABLE 11-2.

RAE Data Tapes Processed Microfilm Read Frequency

13 Jan 69-03 Jul 69 26 Feb 69-12 May 69 450-6550
06 Feb 70-07 Jul 70 06 Mar 70-09 Jun 70 1300;6550
01 May 71-12 Aug 71 1 May 71-12 Aug 71 6550

Processing of the data tapes was performed via the IBM 360/91

and the Univac 1108 computers at NASA-GSFC, Greenbelt, Maryland.

An outline of the microfilm production scheme is illustrated in

Figure 11-8. Each magnetic tape contained 16 days of RAE-1

data in digital form. The data on the tapes were organized into

units known as data sets, each set containing twenty-three words.

Besides upper-V and lower-V antenna temperatures, each data set

also contained the instantaneous satellite pointing direction

in right ascension and declination, impedance probe information,

thermal probe information for calibration purposes, and elapsed

time in msec from 1 Jan 68. All of this information, excepting

the impedance and calibration probe data, was necessary for the

generation of the microfilm plots. Only the Ryle-Vonberg

coarse signals (word numbers 12, 13, 15, and 16 on each data

set) were used.

A Fortran IV program (MPLOT) was written to control the pro-

cessing of the microfilm on a production basis. MPLOT ordered

the data such that individual satellite orbits were displayed

on each frame of the microfilm. In addition, knowledge of the



IBM 360/91

Iplotter I


SD4060 CRT Plotter

Fortran program: Mplot
reads RAE-I tape
calculates orbilal phase
instructs SD4060 platter

16mm Microfilm Output

FIGURE 11-8. Schematic Outline of the Microfilm Processing. The
facilities of NASA-GSFC were used throughout. The
SD (Stromberg Datagraphix) 4060 CRT plotter is a
fast-plotting accessory of the'Univac 1108 computer.

position of Jupiter and of the instantaneous pointing direction

of the upper-V antenna permitted the calculation of the instant

of Jupiter beam transit during each orbit. Based upon this cal-

culation, MPLOT adjusted the time of beam transit (defined as

the 00 orbital phase point) such that it occurred in the middle

of each frame. For ease in reading the microfilm, however, the

orbital phase was shifted by 1000 so that negative numbers

would not appear. As the probability of positively identifying,

or even detecting, Jovian emission decreases rapidly as the

angle from the E-plane beam axis increases, only 56% (2000) of

each orbit was actually plotted. Each frame of the resultant

microfilm presents data at a single RV frequency. An example of

two such frames is illustrated in Figure 11-9.

Two Consecutive Frames of the 16-mm Microfilm Are Shown. Each frame plots the
upper-V and lower-V antenna temperatures as a function of satellite orbital
phase at a single RV frequency. The instant of time coinciding with 00 orbital
phase appears on each frame in the format YYMMDD and HHMM at the top and bottom
respectively. Jupiter beam transit occurs at 1000 orbital phase. An intense
Jupiter storm is apparent at 3930 kHz between 1000 and 1060 orbital phase.

FIGURE 11-9.

11 Ill
II11 A Bll

I I ll ~IE MEll I
*l 111111

I I ll I I
IIN 11 5 In1
Is lls Ill

IN 11 I l


ill I in
I1 I t II I
No 1is 111
II a ll



One ring to rule them all,
One ring to find them,
One ring to bring them all,
And in the Darkness bind them . .
--J. R. R. Tolkien

IN EVALUATING THE PERFORMANCE of an antenna system, there are

two major considerations: (1) The alignment of the antenna's

main beam must be known to within a suitable fraction of the

beamwidth depending upon user requirements; (2) The absolute

gain of the antenna and its variations with azimuth and altitude

must be known in order to make quantitative measurements of

radio source flux densities.

The determination of beam alignment and absolute gain becomes

far more critical and time consuming as the antenna system in-

creases in size. Small antennas, such as the Yagi, lend them-

selves to suitably accurate beam alignment through simple

bore-sighting. Absolute gain may be determined by way of calcu-

lation alone. Unfortunately, these generally straightforward

techniques become more open to question as the system complexity

increases. However, the greater directivity of a large antenna

system makes possible the detection of radio sources whose

positions and absolute intensities are sufficiently well known

that they may be used to map the antenna pattern precisely and

to determine its absolute gain. The technique of calibration by

means of standard radio sources has been chosen for evaluating

the performance of the 26.3 MHz array.

The detection system used in making the observations at 26.3

MHz is of fundamental design. The signals were recorded in the

total-power mode by means of single-conversion superheterodyne

receivers. Each receiver (one for each of three channels) has an

IF bandpass of 250 kHz, with a diode detector and a postdetec-

tion time constant of 1 sec. The receivers were designed and

constructed by Mr. Jorge Levy and Mr. Richard Flagg. Three

different models of chart recorders were used (not simultaneously)

at various times during development of the system to record data:

a dual-channel (PDP) Texas Instruments chart recorder, a

Honeywell three-channel Electronik 16, and a Rikadenki three-

channel chart recorder. The latter instrument was found to be

the most versatile and reliable. The scintillation events for

the phase calibration of the array and the radio source drift

scans for the gain calibration were recorded at 20 cm/hr and

6 cm/hr respectively. The calibration methods are described in

the following sections.


The beam alignment of the array is determined by a complex

network of electronically and manually inserted phase delay

lines between its 640 dipole elements. Beam alignment is

therefore equivalent to phasing and will be referred to as

such. The following is a presentation of the results of the

phase calibration of the array.

Upon completion of the north half of the antenna (320 ele-

ments) a complete phasing unit was established. That is, a test

of the resultant beam positioning represented an evaluation of

the fine (Butler matrix) and coarse (plug-in phasing cables)

phase adjustments which would be incorporated into the entire

array. The Butler matrices and the manually inserted phasing

cables would be required to be working properly in order for

the tests to be successful. During the tests made at this time,

there were five Butler matrices controlling east-west (ew)

phasing and one controlling north-south (ns) phasing.

The method used for testing the ew phasing accuracy was

simply to allow a strong radio source to drift through each

of the eight ew antenna beams in turn. Because the radio

sources chosen for these tests were effectively point sources

relative to the 60 ew HPBW, the resultant drift scans were a

true representation of the actual power pattern of the array,

after the application of receiver calibration data. The normal-

ized field pattern for a linear (ew) array of n dipoles is given

by Kraus (1950):

sin n',/2
n 1 1 sin P2/2

where n = number of elements in array factor = 8, EB = ew field

pattern of a single half-wave dipole, T1= 8d1 sin 4, T2 = Sd2 sin 4,

8 = 2a/f, dI = 0.6 1, d2 = 1.2 1, # = angle (ew) from beam transit

plane. An additional factor contributing to the field pattern,

sin (7/2 cos 4), is due to the combined effect of a reflecting

ground screen and the surface of the ground, each located 1/4

below the dipoles. The variation in this term is insignificant

over the extent of the main beam, however, and has been neglected


In Figure III-1 are shown the results of two such drift-scan

analyses using 3C274 (Virgo A) and 3C218 (Hydra A) taken during

times when there was little ionospheric scintillation to mar the

records. It is clear that the observed drift scans conform very

well to the shape of the theoretically expected power pattern.

Additionally, the expected time of beam transit is correct to

within +1 min. Visual inspection of similar drift scans through

the remaining beam positions led to the same conclusion: the ew

phasing is quite accurate.

Testing the phasing accuracy in the ns direction is somewhat

more complicated. The method utilized takes advantage of the

fact that there are three slightly overlapping ns beams during

normal operation of the array, corresponding to the three

transit transit


Comparison of the Theoretical, Normalized ew
Power Pattern of the Array (Solid Curve) with
the Experimentally Derived Signal Powers
(Solid Circles).


recording channels. The three beams were separated by 30 and

intersected at the -1.5 db point when only half of the array

was used (with the full array, these quantities are different).

Hence in general, a source will appear simultaneously in two of

the beams during a drift scan, with the signal power in each

beam being determined by the position of the two beams relative

to the source. Assuming that the position of the beams is known

precisely and by comparing the deflections occurring simul-

taneously in the two channels, one may determine the position

in declination of an unknown source. Conversely, if the position

of the source is already known, one may use the experimentally

derived power ratios to determine the error in the beam alignment

of the antenna. The latter approach has been used to establish

the ns phasing accuracy of each of the eight ew beams by com-

paring the true position of the source with the position as

deduced from the multiple-beam method. The normalized field

pattern profile of the beams in the ns plane is given by

1 sin n1 /2 sin n2 2/2
nE n2 sin '1/2 sin T2/2 ( -2)

where n1 = 4 and n2 = 5. 41 and 2 are defined as in Equation

III-1 with = angle (ns) from beam axis, d1 = 0.60 X, and

d = 2.40 A. Subscript-1 quantities refer to a ns row of four

dipoles in a group of eight, and subscript-2 quantities to the

array factor formed by the center of the five groups of sixty-

four dipoles each. It was found that a strongly scintillating

source was the best suited for this type of analysis.

The results of analyzing 225 scintillation "events" using

3C144 (Taurus A) are shown in Figure III-2. Each scintillation

event lasts for about 30 sec to a minute, and on the average,

about ten events during a given 25-min drift scan are of

sufficient magnitude for scaling. In addition, a limited number

of events (15) were studied using the radio source 3C218

(Hydra A). This was done to uncover possible variations in

phasing accuracy with zenith angle. The zenith angles at

meridian transit for 3C144 and 3C218 were 80 south and 420

south respectively.

The ns phasing error for each ew beam position may be ob-

tained from Figure III-2. The dotted line represents the true

declination of 3C144 and 3C218. The plotted points represent

the positions of the sources as deduced from the foregoing pro-

cedure. The mean phasing error is 0.34 north, with the error

tending to increase away from meridian transit. The maximum

deviation (4W beam) is 0060 north. Although the data sampled at

-12 declination (3C218) are incomplete, there does not appear

to be any indication of increased phasing error at negative

declinations. Ionospheric refraction, which is expected to con-

tribute insignificantly to the phasing error at 26.3 MHz

(Viner 1975), has been neglected.


- 3C 144

0 3C 218

----- -------
-e Q


4E'3E'2E' IE' IW'2W'3W'4W






r'3O C)



The Declinations of 3C144 and 3C218 As Deduced
from the Method Described in the Text Are
Plotted for Each ew Beam Position. The dotted
line represents the true declination of the
sources. The ns phasing error for each ew
beam position is given by the difference
between the true and calculated declination.





In view of the results presented here, one may expect to be

able to point the main beam of the array to within a probable

error of about 0 6 in the ns plane. This is in fact the angle

between the ns beam positions when the entire array is operating.

From a practical standpoint this means that the values tabulated

in the PTP booklet (Desch 1972) will be correct to within + one

digit in the SNS number. Extensive monitoring of radio sources

over a wide range of zenith angles for the purpose of gain

calibration (see next section) has further substantiated this

result. In addition, the monitoring of very weak Jupiter bursts

over a period of 3 years has been eminently successful (Desch

et al. 1975). It was found that the radio bursts could be dis-

tinguished from interference in three-channel recordings (one

channel for each ns beam) based solely upon the signal power

ratios predicted from the theoretically expected beam positions.

The system phasing appears to be stable over time scales of at

least several years.


The absolute gain of the array has been established over

approximately 600 of zenith angle through drift scan analyses

of thirteen radio sources. Table III-1 tabulates the sources

chosen along with their positions, flux densities in Jy at

26.3 MHz, alternate names, and specific comments concerning

certain of the sources which posed special problems. The flux

TABLE III-. Standard Calibration Sources.

R.A. Dec Flux Density
Source 1950 1950 26.3 MHz
Numbera Source Name h m o Jy Alternate Names Commentsb
(1) (2) (3) (4) (5) (6) (7)

1 PKS 0213-13.2 02 13 -13 13 110+ 10 MSH 02-105 a
2 MSH 02-07 02 18 -02 12 120+ 20 3C 63 b
3 MSH 03-301 03 19 -37 18 3,100+ 300 Fornax A, PKS 0320-37 c
4 3C 89 03 32 -01 23 160+ 30 MSH 03-03, PKS 0331-01
5 3C 144 05 31 21 58 2,500+ 100 Taurus A
6 3C 181 07 25 14 45 60+ 10 PKS 0725+14 d
7 3C 190 07 59 14 27 63+ 10 PKS 0748+14
8 MSH 08-404 08 20 -42 48 1,650+ 100 Puppis A e
9 3C 218 09 16 -11 53 1,800+ 150 Hydra A, MSH 09-104
10 3C 225 09 39 14 17 140+ 30 PKS 0939+14
11 3C 274 12 28 12 40 4,500+ 300 Virgo A
12 MSH 13-402 13 22 -42 46 20,000+3,000 Centaurus A f
13 MSH 22-17 22 12 -17 17 400+ 100 PKS 2211-17, 3C 444

aSource numbers are keyed to Figure III-8.

bKey to comments: (a) confusing source PKS 0213-13.5, (b) confusing source MSH 02-012; (c) extended
source, integrated flux is tabulated, (d) confusing source 4C+15.2, (e) confusing source MSH 08-405,
(f) extended source, integrated flux tabulated.

densities and uncertainties quoted have been obtained from the

spectra presented in Figures III-3 through III-7. In each case

the uncertainty in the flux density at 26.3 MHz is arrived at

after making a reasonable estimate of the variation in the flux

density curve because of the scatter in the published points.

This scatter is seen to be greatest at the lower frequencies.

The higher frequency data points (>26 MHz) have been obtained

from figures published by Dixon (1970). The flux densities at

decameter wavelengths were obtained from various sources,

namely Erickson and Cronyn (1965) at 26.3 MHz; Roger et al.

(1969) at 22.25 MHz; Bridle and Purton (1968) at 10.03 MHz;

Braude et al. (1969) at 12.6, 14.7, 16.7, 20.0, and 25.0 MHz;

Bazelyan et al. (1965) at 20.0, 25.0, 34.0, and 38.5 MHz;

Guidice (1966) at 22.3, 26.7, 33.45, and 38.7 MIz; and Shain

(1958) at 19.7 MHz.

Once the standard radio sources have been chosen, the pro-

cedure for establishing antenna gain is relatively straight-

forward. In Table III-2 are presented the results of the drift-

scan analyses, along with the dates that the records were taken

and the phasing parameters used in locating the sources. The

following are the steps used in making a single determination of

the antenna gain at a particular zenith angle.

1. A drift scan is made of one of the sources
tabulated in Table III-1 using the phasing
parameters appearing in Columns 3 through 6
of Table III-2.

I I I I1111 A I I1111 I I I il


0 C 181





.L 10



\ 0

I I I I 0.1



FIGURE 111-3. Flux Density Spectra of 3C225 and 3C181.
See text for references.

. I 1 1 I I I I l l


. I


I I I I 11111






10 102


FIGURE 111-4. Flux Density Spectra of PKS
See text for references.


0213-13.2 and 3C190.

U I _- i i 11 ill -- I I 1111111 J I I 11 iit

1 PKS 0213-13.2



10 0 -0




II I .I t ll! I I I Iuliu l IoI I IIIIl


FIGURE III-5. Flux Density Spectra of MSH 22-17, MSH 02-07, and 3C89.
See text for references.


S0 MSH 03-301
10 0o A 3C 218 -


X 0,,


10 \ o o

i0 1 i. I .. I I . I l l i I I I i I I 1.l,

10 102 103 10


FIGURE III-6. Flux Density Spectra of 3C274,
See text for references.

MSH 03-301, and 3C218.



O -

10 102 103 104


FIGURE 111-7. Flux Density Spectra of MSH 13-402, MSH 08-404, and 3C144.
See text for references.

TABLE 111-2. Results of Array Gain Calibration

T T* A*
Phasing Parameters TA TA e
Source Date CNS FNS RSP EW P P2 (K) (K) (K/Jy)
(1) (2) (3) (4) (5) (6) (7) (8) (9) (10) (11)

PKS 0213-13.2 17 Jan 75 +3 5 2 1E 0.69 1.39 17+ 3 18+ 2 0.16+.02

1.39 22+ 4
1.39 24+ 7
1.23 28+ 9
0.79 119+11
36+ 6
36+ 4
29+ 5
22+ 4
30+ 5
28+ 8
1.49 33+ 5
1.49 26+ 5
19+ 7
14+ 4
22+ 4
1.82 165+ 5
1.82 205+ 5
273+ 9
31+ 5

MSH 02-07
MSH 03-301
3C 89

3C 144

3C 181

3C 190

MSH 08-404

3C 218

3C 225


23+ 4
25+ 6
59+ 7
47+ 8
36+ 7
49+ 8
29+ 5
23+ 5
21+ 6
32+ 6
156+ 5
190+ 5
45+ 7


TABLE 111-2, continued:

T T* A*
Phasing Parameters A A e
Source Date CNS FNS RSP EW P1 P2 (K) (K) (K/Jy)
(1) (2) (3) (4) (5) (6) (7) (8) (9) (10) (11)

3C 274 26 Feb 74 +1 6 5 1W 0.86 1,277+78 1,480+90 0.33+.03
22 Apr 74 +1 6 5 1E 0.86 1,287+75 1,491+87 0.33+.03
MSH 13-402 14 Apr 73 +4 7 5 1W 0.69 0.45 404+24 1,309+78 0.07+.01
15 Apr 73 +4 7 5 1W 0.69 0.45 338+19 1,096+62 0.05+.01
17 Apr 73 +4 7 5 1E 0.69 0.45 406+ 6 1,316+20 0.07+.01
20 Apr 73 +4 7 4 1E 0.70 0.45 425+ 6 1,344+18 0.07+.01
21 Apr 73 +4 7 4 1W 0.70 0.45 494+17 1,564+53 0.01+.01
MSH 22-17 10 Jul 73 +4 2 7 1W 0.61 30+ 9 45+15 0.12+.05

NOTE: Key to Table 11-2: (1) name of calibration source, (2) date of observation, (3-5) phase-
tracking parameters used in making observations, as derived from PTP booklet, (6) EW beam through
which drift scan occurred, (7) normalized power pattern of antenna in direction of source,
applied to TA as a correction, (8) correction applied to TA due to presence of confusing source,
or due to a source having a finite extent relative to main beam; blank column--no correction
necessary, (9) measured, uncorrected antenna temperature due to source alone, (10) corrected
antenna temperature, (11) final effective area of antenna from Equation III-5.

2. The antenna temperature Tn due to the combined
effects of the galactic background and the
source is calibrated against a secondary noise
standard at 5- to 10-min intervals as detailed
in Appendix A.

3. Usable records (%25%) are those which are rela-
tively free of large antenna temperature fluctua-
tions caused by ionosopheric scintillation. The
remaining 75% are not used. Occasionally the
calibration will exhibit a constant drift. That
is, because of ambient temperature changes at
the field preamplifier (see Appendix A), a given
calibration level will record at different levels
on the chart recorder. This thermal drift is al-
ways subtracted out of each drift scan before

4. The antenna Temperature TA due to the source
alone is determined by subtracting the galactic
background from the temperature due to the source
and background combined. The temperatures used
are those which occur at beam transit when the
source deflection is greatest. The galactic
background temperature at beam transit is ob-
tained by drawing a straight line between the
values for the antenna temperature obtained at
the antenna pattern nulls which occur approxi-
mately 25 min before and after beam transit (see
Figure III-1).

5a. This value for TA must then be corrected by a
factor p which is the value of the antenna's
normalized power pattern in the direction of the
source. This factor is always less than 1.0 be-
cause in general a source will never pass through
the beam axes of each of the modulating power
pattern envelopes of the array simultaneously.
In all, three ns envelopes and two ew envelopes
combine as factors to generate a final product
p which is different for each source and set
of phasing parameters used. In calculating p,,
Equations III-1, III-2, and the principle of
pattern multiplication are used (Kraus 1950).
The corrected figure TA = TA/pI is thus the value
for the antenna temperature normalized to beam
axis conditions.

5b. In the case of four of the radio sources (PKS
0213-13.2, MSH 02-07, 3C181, MSH 08-404), an
additional correction P2 is necessary because
of the presence of confusing sources as noted in
Column 7 of Table III-1. The confusing sources
were close enough to the standard radio sources
that their contributions to the antenna tempera-
tures were not negligible. In these four special
cases the final corrected figure TX is given by
TA/P1 P2

5C. A further correction was applied to two of the
sources (MSHI 03-301, MSH 13-402) due to their
finite extent relative to the main beam of the
antenna. This correction is necessary if the
angular diameter of the source exceeds about one
fifth of the antenna half-power beamwidth (Gui-
dice and Castelli 1968). At decameter wave-
lengths the ns by ew extents of MSH 03-301 and
MSH 13-402 are approximately 102 by 1 75 and 4?5
by 2 0 respectively (Shain 1958). Given the ns
and ew antenna half-power beamwidths, one may
then calculate the correction factor to be used
for each source (Baars et al. 1965). The values
so derived are tabulated under p2 in Column 8 of
Table III-2.

6. The absolute gain of the antenna is now directly
expressable in terms of quantities which are
tabulated in Tables III-1 and III-2. From Kraus
(1966) the effective area of an antenna is given

2kT 1026
A (111-3)
e S

where Ae = effective area in m k = Boltzman's
constant = 1.38 x 10-23joules/K, S = true flux
density of source in Jy, TA = antenna tempera-
ture K due to source of flux density S,= T .
It has been decided to express the absolute gain
in terms of K/Jy which are the most practical
dimensions, both from the standpoint of quantities
derived by the foregoing procedure and applica-
tion to future observations. Collecting constants,
one can see Equation 111-3

A = 2.76 103 A* (III-4)
e e

where A* = effective area of antenna in working
units of K/Jy. As mentioned, A* may be obtained
directly from quantities already derived in the
gain calibration procedure, namely the flux
density S and the adjusted antenna temperature
T*. That is, A* is given by

A* = T*/S (111-5)

This quantity is tabulated for each of the drift
scans listed in Table III-2.

The errors quoted in A* are due to two effects. The uncertainty

in scaling the chart recordings, primarily as a result of scin-

tillation effects, contributed about 50% to the final error.

Another 50% is due to the uncertainty in the assumed flux density

values as mentioned previously. Uncertainties because of phasing

errors (calculation of p ), calibration of the secondary noise

standard (Appendix A), and changes in the calibrator line attenua-

tion with ambient temperature fluctuations have been neglected.

These rather small effects are not expected to contribute more

than 10 to 15% to the errors already quoted in A*. Ionospheric

absorption, which is rarely greater than 2% for nighttime propa-

gation conditions at 26.3 MHz (Viner 1975), has also been

neglected. All of the observations reported on in this chapter

were made between 1800 and 0600 local time.

In Figure III-8 are plotted the weighted mean values of A* as

a function of zenith angle and declination. There is a general

trend in that the absolute gain of the antenna is seen to de-

crease with increasing zenith angle z, as expected. The data

points do not, however, obey a simple cos (z) sin [a/2 cos (z)]

fall-off as might be expected from the combined effects of

foreshortening (cosine factor) and reduced ground reflection

(sine factor). This may be due to variations in impedance match-

ing with phase which have not been evaluated. In addition, the

value obtained for A* using 3C144 appears to be too low, since
one would expect the gain to peak at 00 zenith angle. In any

event, a greater number of measurements are needed at declina-

tions exceeding 100.

The procedure used here in deriving the absolute gain of the

antenna is considered to be a compromise between that necessary

for precise monitoring of weak, unknown radio sources and that

necessary for determining the flux density of Jupiter radio

bursts, for which precise calibration is unnecessary. The pro-

cedure is well suited for determining relative changes in array

performance over long periods of time in order to diagnose the

possible damaging effects of lightning or the gradual deteriora-

tion of electrical elements with age. Conversely, monitoring

radio sources for intensity fluctuations on time scales greater

than a few days is also possible, especially if a nearby com-

parison source is used. The reproduceability of drift scans

80 70 60 50 40 30 20 10





-500 -400 -300 -200

-100 00

100 200 300


FIGURE 111-8.

The Effective Area of the Array in K/Jy Is
Plotted As a Function of Source Declination
and Zenith Angle. The source numbers are keyed
to Table III-1.


11 4 10 51
~- I II

-08 13T
12* #3


from standard sources over a period of about 2 years indicates

that about a 10% variation in collecting area or source

intensity could be recognized following the procedures outlined



One man's magic is another
man's engineering.
--Robert Heinlein


prior to this study, it was considered appropriate that the

V-antenna power pattern and its attitude relative to the local

gravity vector (i.e. the nominal upper-V pointing direction) be

investigated as carefully as possible prior to undertaking an

intensive campaign. Using the sun as a source of powerful

Type-III bursts, a method was devised to verify the pointing

accuracy in the E-plane and subsequently to confirm the exis-

tence of low-frequency Jovian activity. Further, by superposing

the Jovian activity recorded over many orbital scans, it was

possible to synthesize E-plane antenna patterns at a number of

the higher RV frequencies and thus complement similarly

directed theoretical studies. This preliminary investigation is

described in this section.


The orbital superposition technique is applicable only at

frequencies exceeding 1310 kHz. Above this frequency, the

probability of detecting emission from the sun or Jupiter is a

clearly defined function of satellite orbital phase, after the

easily recognizable terrestrial interference has been deleted

(below this frequency the terrestrial interference is excessive).

That is, as is evident from Figure 11-7, there is a direct rela-

tionship between orbital phase and the pointing direction of

the upper V in the E plane relative to a given source direc-

tion. Consequently, as the upper-V antenna is swept across a

source the resulting antenna temperature will be directly pro-

portional to the E-plane antenna power pattern or equivalently

to the satellite's orbital phase. Clearly, the orbital phase

variation will become more evident at the higher frequencies as

the E-plane power patterns become increasingly directive. Be-

cause of the precession of the orbital plane of the satellite,

the precise response function is also modulated somewhat by the

H-plane pattern, but over relatively short periods of time

(< 3 months) this effect is negligible.

In using a source whose activity is unpredictable, such as

the sun, the data collected over many orbital scans must be

superposed until a smooth antenna pattern begins to emerge. In

practice, the entire observing period in 1969 (26 Feb-12 May)

was used in confirming both the upper-V attitude and the

existence of the Jovian emission.

In order to insure that this initial study be as objective

as possible, the microfilm versions of the data were temporarily

abandoned and the computer was instructed to search the original

RAE-1 magnetic tapes for "upper-V only" events, that is those

which were supposedly of celestial origin. The program simul-

taneously compared the upper-V and lower-V radiometer channels

at 3930 and 4700 kHz, with "upper-V only" events being counted

when the signal from the upper V was found to be at least 3 db

greater than the lower-V signal. Superposing 500 orbital epochs,

the probability of occurrence of "upper-V only" events was

plotted as a function of the orbital phase in Figure IV-1. Here

the 00 orbital-phase point in each orbit occurred at the time

of Jupiter beam transit. As Jupiter was at or close to opposition

during the period, the sun was always near 1800 orbital phase.

As Figure IV-1 clearly indicates, the two significant peaks

at both frequencies occurred at 00 and 1800 orbital phase,

corresponding respectively to the positions of Jupiter and the

sun as seen by the upper-V antenna. This result firmly

established not only the existence of the Jovian activity, but

also the pointing accuracy of the upper-V antenna in the E plane.

Not all data were as well behaved as the 3930- and 4700-

kHz radiometer channels. At 2,200 kHz, erratic drifts of the

upper-V or lower-V background levels would sometimes result in

the computer counting as "upper-V only" events those which by

visual inspection were obviously not Jovian. At 6550 kHz the


-60o 00 600 1200 1800 -120o


Occurrence Probability of "Upper-V Only"
Events As a Function of RAE-1 Orbital
Phase. Contributions from'both Jupiter
and the sun are clearly visible. The
theoretical (500 apex angle) calcula-
tion of Sayre (1974) is included for
comparison at 3930 kHz.

common occurrence of terrestrial ground interference ("ionos-

pheric breakthrough") often resulted in similar misidentifica-

tions. Hence, at these two frequencies the microfilm versions

of the data were examined with the resulting phase variations

appearing quite similar to those found at 3930 and 4700 kHz.

At 6550 kHz, however, only the data within about +500 of

Jupiter beam transit were read (see Figure IV-2).

Besides verifying the existence of celestial activity at 00

orbital phase, it was also possible at 3930 and 6650 kHz to

compare the experimentally derived antenna patterns with the

theoretical calculations for E-plane response made by Sayre

(1974). In both cases, the theoretical pattern shown is that

for the ideal in-orbit case; that is, the V antenna is assumed

to have a 500 apex angle. At 6550 kHz the agreement is

remarkably good, especially considering that one is plotting

the occurrence of sporadic events over a 3-month interval.

At 3930 kHz, the 500 apex-angle profile appears somewhat

narrower than the pattern derived from the orbital superposi-

tion of activity, suggesting that the 350 apex-angle calculation

of Sayre, which results in a less directional main beam, might

be more applicable. The data do not permit a definitive judg-

ment, however.

The first sidelobe, at +600 orbital phase, is also evident at

3930 kHz, although it is displaced somewhat relative to that

predicted by the 500 apex-angle model. The sidelobe at -600

6550 kHz

--- SAYRE (1974)



0.3 -

0.0 1-

'L- -

-40o -300 -200 -IO

100 200 300 40


Normalized Occurrence Probability of "Upper-V Only"
Events As a Function of Orbital Phase at 6550 kHz.
The theoretical (500 apex angle) calculation of
Sayre (1974) for the E-plane antenna power pattern
is included for comparison.




- - -

orbital phase does not appear at all, probably because of the

masking effect of the ionospheric breakthrough which frequently

occurred over this range of orbital phase.

Further comparison of the patterns in Figures IV-1 and IV-2

reveals a broadening with decreasing RV frequency. This is to

be expected, of course, because of the decrease in directivity

of the V antenna with increasing wavelength. The antenna

patterns at 180 phase are also broader than the 0 -phase pro-

files occurring at the same frequency. This is due to the fact

that the sun did not remain precisely at 1800 with respect to

Jupiter during the 3-month interval under study.

The results presented in this section have demonstrated the

usefulness of both computer reading of the RAE-1 tapes and

visual examination of the microfilm transcriptions in identifying

Jovian activity. The remainder of the data has been reduced via

the microfilm-inspection method for reasons explained in Section

11-4. In the following section, some of the results obtained in

the foregoing analysis are used to assist in deriving values

for the absolute gain of the V antenna at frequencies between

450 and 6550 kHz.


The calibration method employed at 26.3 MHz, which depended

upon the detection of standard radio sources, was not applicable

in the case of the RAE system. No continuously emitting "point"

sources have been identified using the RAE-1 satellite, even

at the higher frequencies where the antenna directivity is

greatest (Alexander, personal communication). As a result, the

V-antenna effective areas were determined by an indirect method

(Kraus 1966) which is outlined below. The method requires a

knowledge of the antenna's power pattern and also a knowledge of

the loss factor involved in the matching of the antenna to the

RV receiver. The former is known to some extent both theo-

retically and experimentally as discussed in the previous

section. The latter depends upon the absolute measurement of the

temperature of some sky source--thus permitting the direct

comparison of V-antenna-derived source temperatures with the

"true" source temperatures. In the present analysis, the true

source temperatures have been determined from the absolute

brightness measurements of the galactic background made by

Brown (1973) and by Alexander et al. (1969). They are tabulated

in Column 2 of Table IV-1 for each of the RAE frequencies. Hence

the loss factor k is established at each frequency by forming

the ratio of T (background temperature measured by the V-antenna)

over T (known brightness temperature of the galactic background).

These figures are tabulated in Columns 3 through 5 of Table IV-1.

The absolute brightness measurements of the galactic back-

ground are critically dependent upon an understanding of the

radiation characteristics of the dipole antennas employed by

the above authors. Fortunately, studies of the dipole antenna

TABLE IV-1. V-Antenna Calibration Parameters.

Freq B T T k EM 0HP A*
-20 6 6
kHz 10 10 10
(1) (2) (3) (4) (5) (6) (7) (8) (9)

6550 1.20+18% 0.91 0.4 +10% 0.44 +20% 0.72 14 +10% 60 +25% 2600 +33%
4700 1.22+18% 1.80 1.0 +10% 0.55 +20% 0.73 30 +20% 62 +25% 2900 +38%
3930 1.30+18% 2.70 1.0 +10% 0.37 T20% 0.80 38 +20% 64 +25% 2300 +38%
2200 1.10+12% 7.40 4.0 +10% 0.54 +16% 0.95 60 +25% 100 +25% 5200 +39%
1310 0.73+10% 14.00 4.0a+10% 0.29a+14% 0.95 180 +25% 120 +25% 2200a+38%
900 0.49+10% 20.00 4.0a+10% 0.20a+14% --b -.b --b 2600a+14%
700 0.35+11% 23.00 2.0a+10% 0.09a+15% --b .--b --bT 2000a+15%
450 0.10+14% 16.00 1.0a+10% 0.06a+17% --b -__bT -bT 3200a+17%

aOnly nominal values are tabulated (see text).

bAntenna pattern effectively that of a short dipole.

NOTE: Key to Table IV-1: (1) RAE-1 frequency in kHz, (2) galactic background brightness in
units of Wm-2Hz- sr-1, (3) brightness temperature (K) from (2), (4) brightness temperature
(K) from RAE-1, V-antenna measurements, (5) loss factor dimensionlesss), (6) main-beam
efficiency, (7) E-plane HPBW in degrees, (8) H-plane HPBW in degrees, (9) V-antenna
effective area in m2.

have been quite exhaustive, and its behavior in the plasma

environment of the earth's magnetosphere is well documented.

Nonetheless, it is well to keep in mind that the method

employed here to calibrate the V-antenna gain is not as

reliable as that used in calibrating the effective area of the

26.3 MHz array, primarily because of its ultimate dependence

upon measurements made by a secondary antenna system.

The following is a derivation of the equations used in

determining the effective area of the V-antenna at a given

frequency. Kraus (1966) has shown that

A A = A/ (IV-l)

where Ae = antenna effective area as determined solely from

the antenna pattern and 2A = beam solid angle in sr (sr = stera-

dian = radian2). In the ideal case, when ohmic and mismatch

losses are negligible, Equation IV-1 is correct as stated.

However, in general an ohmic loss factor k must be included.

Equation IV-1 then becomes

A* = k 2 /lA (IV-2)

where A* = true effective area, k = loss factor because of im-
pedance mismatching, resistance, etc., and in general,
pedance mismatching, resistance, etc., and in general, 0 < k <1.0.

In the present analysis, k. is given by the ratio T /T as ex-

plained previously. The beam solid angle may be expressed as

"A = "M/M (IV-3)

where QM = main beam solid angle and E = main beam efficiency.

Furthermore, the main beam solid angle may be expressed as

QM = kp EHP HP (IV-4)

where k = pattern shape factor = 1.05 + .05, OHP = H'BW in

radians in the 0 direction (E plane), and ,P= HPBW in radians

in the ( direction (H plane). Because k varies over so narrow
a range as compared to the relative uncertainty involved in

determining the other factors in this analysis, we shall set k
equal to 1.0. Then, by combining the above equations and con-

verting from radians to degrees and from wavelength to frequency,

we obtain the final form:

29.5 x 10 E k(IV-5)
A* -
e 0 0 2

where oi H = E-plane HPBW in degrees, oHP = H-plane HPBW in

degrees, v = frequency in MHz, and A* = true effective area in m

Equation IV-5 has been used to calculate the effective area

of the V-antenna at frequencies between 6550 and 1310 kHz

inclusive. At the lower frequencies the V-antenna power pattern

is essentially that of a short dipole. Hence, the effective

area at 900, 700, and 450 kHz has been calculated from the


A* = 0.119k X (IV-6)
e 0

which is the equation for evaluating the effective area of a

short dipole (Kraus 1950). Using Equations IV-5 and IV-6 we

have tabulated the true effective areas in Column 9 of Table

IV-1. The remainder of this section will be devoted to de-

scribing how the parameters used in evaluating A* and appearing

in Table IV-1 were obtained.

The brightness spectrum of the galactic background derived

by Brown (1973) has been used for the frequencies between 2200

and 450 kHz because the uncertainties quoted were less than

those of Alexander et al. (1969). Hence, the error limits

appearing in Table IV-1 at these frequencies were those obtained

by Brown. His investigation, which employed the 91-m dipole

on the IMP-6, did not extend above 2600 kHz however. Thus,

between 6550 and 3930 kHz, the figures derived by Alexander

et al. have been used. Using the 37-m dipole on the RAE-1,

Alexander et al. have formally calculated an error in each

brightness measurement of +18%. They estimated that the abso-

lute determination of their brightness values might have been

in doubt by as much as +25%, however, because of systematic or

random errors not accounted for. The formal error (+18%) has

been chosen here as representative of the uncertainty since

the lower frequency extension of the RAE-1 curve merges so well

with the Imp-6 spectrum obtained by Brown. The brightness

temperatures were obtained directly from the brightness measure-

ments themselves by means of the equation

T = 0.5A2 B/k (IV-7)

where B = galactic background brightness (Wm-2 Hz sr-) and

k = Boltzmanifs constant. The percentage errors associated with

T (not explicitly shown) are thus the same as those associated

with B.

The galactic background temperatures T measured with the V-

antenna system on the RAE-1 were read directly from the micro-

film plots of the upper-V antenna temperatures (see Figure 11-9).

They are subject to an uncertainty of about +10% at each RV

frequency. This value has been derived by Alexander and Novaco

(1974), based upon an analysis of the errors inherent in the

RV-radiometer measurements. The individual factors contributing

to the total uncertainty were (i) receiver noise, (ii) internal

satellite temperature uncertainties, (iii) finite digitization

steps involved in telemetry, (iv) errors in the radiometer

calibration, (v) gain shifts in the radiometer, and (vi) varia-

tions in antenna impedance because of in situ plasma density

changes and V-antenna boom motions. Hence, combining the errors

associated with TS and T we obtain the uncertainty in k

which, depending upon frequency, is seen to vary between +14%

and +20%.

Changes in the galactic background level because of Factor

vi are quite often apparent and should be discussed here in

greater detail. The phenomenon manifests itself as a background

level oscillation occurring with a period of about 224 min (one

orbital period). During this time the mean plasma density will

vary by a factor of about three, thus changing the degree to

which the V-antenna is matched to the RV receiver system. This

effect is minimal at frequencies above about 4000 kHz

(Alexander and Novaco 1974). Often superimposed on this drift

is a shorter-period variation in the background temperature

because of changes in solar insolation on the V-antenna booms.

The thermal stress on the booms during times of solar shadowing

establishes a physical oscillation with a period of 50 min

which also changes the effective impedance match. Because of

these effects, only the nominal values of T k and hence, A*

appear in Table IV-1 at frequencies below 2200 kHz. In

general, the variations were observed to decrease with increas-

ing frequency and were not apparent above 1310 kHz.

The main beam efficiencies, tabulated in Column 6 of Table

IV-1, have been estimated from the published theoretical V-

antenna power patterns (Weber et al. 1971; Sayre 1974). The

main-beam efficiency is the fraction of the total power inci-

dent over 47 sr which is intercepted by the antenna's main

beam. Since the numbers are only estimates and do not involve

experimentally derived data, no uncertainties have been

attached to them.

Where possible, HPBW's have been derived from the orbital

superposition data presented in Section IV-1. Hence, the E-plane

values (Column 7, Table IV-1) quoted at 6550, 4700, and 3930

kHz are the experimentally derived figures. It is of interest

to compare these values with the theoretical predictions of

Sayre (1974). As pointed out in the previous section, the

6550-kHz pattern derived from the stacking of Jupiter bursts

agreed exceptionally well with Sayre's analysis for the 500

apex-angle case. The 3930-kHz value is approximately midway

between the 500 and 350 apex-angle calculation of Sayre. No

theoretical 4700-kHz antenna patterns have been calculated as

yet. The H-plane HPBW's at 6550 and 3930 kHz were taken from

Sayre's 350 apex-angle calculation. The average of these two

values was used at 4700 kHz. The greater uncertainty in

choosing the proper H-plane value is reflected in the generally

larger errors associated with them. At 2200 and 1310 kHz the

figures quoted are from Alexander (personal communication).


Below 1310 kHz, power pattern information was not necessary

for determining effective area.

As a final note, it is evident that the uncertainties in

A* are rather large. It should be clear from the discussion,

however, that the figures are not rigorous determinations of

the errors but rather merely representative of the general

degree of certainty which may be attached to each figure.


It has always seemed to me that
the most difficult part of
building a bridge would be the
--Robert Benchley


ducted from 16 April 1973 to the present. The monitoring has

not been carried out entirely without interruption, however,

as observations are normally suspended for a 2- to 3-month

interval centered on Jovian conjunction. Additionally, during

the regular season, the program may be interrupted for several

days at a time for special experiments which require rephasing

(e.g. lunar occultation, VLBI, or QSO monitoring). Unscheduled

interruptions such as those due to severe lightning conditions

or equipment failure have generally been a minor source of

difficulty, although summer atmospheric conditions can reduce

operations by as much as 20 to 25%. Nevertheless, the gross

statistics derived from the uncontaminated Jupiter data (Table

V-l) are impressive, especially in view of the fact that obser-

vations are possible for only about 3.4 out of every 24 hours.

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