SMITHSONIAN

CONTRIBUTIONS

to

ASTROPHYSICS

VOLUME I

SMITHSONIAN

INSTITUTION

WASHINGTON

D. C.

Contents of Volume 1

No. 1. New Horizons in Astronomy. Pages i-x, 1-181. Published December 19, 1956.

No. 2. Papers on Reduction Methods jor Photographic Meteors.

Pages 183-243. Published May 13, 1957.

Smithsonian

Contributions to Astrophysics

VOLUME 1, NUMBER 1

NEW HORIZONS IN ASTRONOMY

edited by FRED L. WHIPPLE

DIRECTOR, ASTROPHYSICAL OBSERVATORY

SMITHSONIAN INSTITUTION

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SMITHSONIAN INSTITUTION

Washington, D. C.

1956
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Smithsonian Contributions to Astrophysics

The Smithsonian Institution's Astrophysical Observatory is passing

through a period of reorganization. The scientific headquarters have

been moved to Cambridge, Mass., where a close working association with

the Harvard College Observatory is being developed. These changes, as

well as the many advances in science and technology, force us to reevaluate

the basic scientific policies and goals of the Astrophysical Observatory.

With the present rapid progress in basic science and its application to

technology, research in the fundamental astrophysical processes of the

sun, earth, planets, and interplanetary medium is going forward at a

rapid pace. Concurrently, our mushrooming technology has become more

and more sensitive to these phenomena of the solar system, heretofore

thought to be of only academic interest.

In keeping with these developments, the Astrophysical Observatory

must broaden the scope of its activities while continuing its long-established tradition of research—a tradition that began with the pioneering

studies on aerodynamics by the observatory's founder, Samuel P. Langley,

and continued during years of active investigations under Dr. C. G.

Abbot in solar and terrestrial phenomena and their interrelationships.

Thus, our future program will include not only radiation, its variations,

and its effects on our planet, but also other phenomena that affect the

earth and its atmosphere. We find, for example, that corpuscular radiation from the sun is associated with solar activity as well as with such

geophysical phenomena as the aurorae, night air glow, terrestrial magnetism, the ionosphere, radio transmission, and other observed phenomena.

Meteoric bodies, dissipated from comets or chipped from asteroidal

sources, not only produce meteors in our atmosphere and provide samples

of cosmic material for our museums but also ionize the high atmosphere

and contribute to its chemical composition, scatter sunlight in the zodiacal

light, and provide us with a cosmic ballistics laboratory of ultravelocity

projectiles.

The Astrophysical Observatory, therefore, faces widening horizons.

It will broaden the scope of its studies of solar phenomena and atmospheric effects to include related phenomena within the solar system such

as meteorites, hypervelocity ballistics, and the zodiacal light. It will

cooperate in exploiting new methods of research and technology such as

nuclear processes and artificial satellites. To strengthen its effort the

Astrophysical Observatory will not carry out its research in isolation;

it will coordinate its work closely with that of other astrophysicists and

geophysicists who are studying similar problems.

Although research in astrophysics continues to expand, no corresponding growth has occurred in the avenues of publication. Observational

results are often so compressed in print that they cannot be analyzed in
detail by independent investigators; and advances in technical procedures

or instrumentations are often presented with such brevity that their value

to the potential user is small.

In an effort to affect these trends and to provide a proper communication for the results of research conducted at the Astrophysical

Observatory, we are inaugurating this new publication, Smithsonian Contributions to Astrophysics. Its purpose is the "increase and diffusion of

knowledge" in the field of astrophysics, with particular emphasis on

problems of the sun, the earth, and the solar system. Its pages will be

open to a limited number of papers b\r other investigators with whom we

have common interests.

We earnestly hope that Smithsonian Contributions to Astrophysics will

serve to promote a greater understanding, appreciation, and enjoyment of

that part of the universe in which we are privileged to live.

FRED L. WHIPPLE, Director

Astrophysical Observatory

Smithsonian Institution

Cambridge, Massachusetts

November 14, 1956
A grant from the National Science Foundation

toward the publication of New Horizons in

Astronomy is gratefully acknowledged.

New Horizons in Astronomy

Preface

The project of collecting a series of papers oriented toward new

horizons in astronomy began in the 1954-55 Panel of Astronomy to the

National Science Foundation. The initial impetus came largely from the

enthusiasm of Dr. Otto Struve, who has presented some of his thoughts

in a paper entitled "The General Needs of Astronomy" (Publ. Astron.

Soc. Pacific, vol. 67, p. 214, 1955). The resulting deliberations led to the

establishment of an ad hoc committee on the "Needs of Astronomy," of

which I was made chairman. I felt that some of the goals of this committee could be attained by publishing a collection of papers by leaders

in the various fields of astronomy who would present their concepts of

research projects most likely to advance the science of astronomy during

the next decade.

The purpose is at least fourfold:

To provide material that might assist the National Science Foundation,

in its allocation of funds, to support the astronomical research most likely

to be productive.

To help younger astronomers choose the most fertile fields for active

research.

To encourage established astronomers to pause for a moment and

reflect on the most advantageous planning in their own fields of research,

and to supply them with a broader insight into the activities in other

fields of astronomy, possibly related to their own.

To encourage research in new or neglected fields of astronomy, particularly those that border on other areas of physical science such as

geophysics, electronics, chemistry, fluid dynamics, and the like.

The publication of New Horizons in Astronomy has been supported

jointly by the Smithsonian Institution and the National Science Foundation. The Astrophysical Observatory shares with the National Science

Foundation the hope that these papers may promote intensified efforts

and increasingly significant results in astronomical research.

The hearty cooperation of the many busy scientists who took time from

their other duties to write these papers should establish a debt of gratitude

in the minds of all those who will profit by these contributions.

I am deeply indebted to Mr. Paul H. Oehser, Mr. Ernest Biebighauser,

and Mr. John S. Lea, of the Editorial and Publications Division of the

Smithsonian Institution, and to Mrs. Lyle Boyd, of the Harvard College

Observatory, for their great interest and effort in this project; otherwise

these papers would not now rest in your hands.

FRED H. WHIPPLE

Contents

Page

Techniques and Instrumentation

Optics: I. S. Bovven 1

Solar Instrumentation: J. W. Evans and R. B. Dunn 5

Astronomical Seeing and Scintillation: Geoffrey Keller 9

Spectroscopy: Charlotte E. Moore 13

Radial Velocities: R. M. Petrie 15

Astrometry: K. Aa. Strand 21

Photoelectric Astronomy: A. E. Whitford 25

Interferometers in Radio Astronomy: B. F. Burke and M. A. Tuve . 31

Related Sciences

Rocketry: R. Tousey 39

Computing Machines in Astronomy: Paul Herget 45

Turbulence: S. Chandrasekhar 47

Astroballistics: Richard N. Thomas 49

Electrodynamics of Fluids and Plasmas: Max Krook 53

Hydromagnetic and Plasma Problems: Eugene N. Parker 59

Solar System

Earth's Magnetism: Walter M. Elsasser 67

Time: William Markowitz 73

Geographical Positions: John A. O'Keefe 75

Celestial Mechanics: G. M. Clemence 79

Meteorites: John S. Rinehart 81

Meteors: Fred L. Whipple 83

Minor Planets: Paul Herget 87

Planets, Satellites, and Comets: Gerard P. Kuiper 89

Airglow and Aurora: A. B. Meinel 95

Solar-Terrestrial Relationships: Weather and Communications:

Walter O r Roberts 99

Solar Physics: Leo Goldberg and Donald H. Menzel 103

Stars and the Galaxy

Spectral Classification: W. W. Morgan . 115

Stellar Atmospheres: Marshal H. Wrubel 119

Aerodynamic Problems in Stellar Atmospheres: Richard N. Thomas. 123

Emission Nebula*: Lawrence H. Aller 127

381014—56 2
CONTENTS

Page

Stars and the Galaxy—Continued

Double and Multiple Stars: Otto Struve 131

Stellar Dynamics: U. Van Wijk 135

Variable Stars: C. Payne-Gaposchkin 137

Interstellar Matter: Lyman Spitzer, Jr 139

21-Centimeter Research: Bart J. Bok 141

Some Observational Problems in Galactic Structure: Harold Weaver. 149

Galactic Clusters and Associations: A. Blaauw 159

Globular Clusters Observed through a Crystal Ball: W. A. Baum. . 165

Stellar Structure and Evolution: M. Schwarzschild 177

Radio Sources: John D. Kraus 179
Techniques and Instrumentation

Optics

By I. S. Bowenl

Geometrical optics is one of the oldest of

sciences. Most of the basic laws have been

known for centuries, yet in the past two or

three decades advances in optical designs have

revolutionized astronomical techniques. An

outstanding example of this is the Schmidt

camera. Until the discovery of the principle

of this instrument in 1931, camera optics

providing critical definition over a field greater

than two or three degrees in diameter were

limited to focal ratios of d/f =1/5 or less.

With the Schmidt camera critical definition

can be achieved over wide fields at focal ratios

as fast as dff=l. This 10- to 100-fold gain in

speed has opened up many new fields in

astronomy.

Furthermore it has been possible to design,

for special purposes, many modifications of the

Schmidt cameras, some of which achieve even

greater gains in speed and field. These include

the Baker Super-Schmidt for meteor photography and the solid-block and thick-mirror

Schmidts, the Schmidt-aplanatic sphere combination, and the twice-through corrector plates

for spectroscopic applications. There are undoubtedly many additional modifications that

should be explored. For example, the focal

surfaces of most of these Schmidt systems are

curved with a radius approximately equal to

the focal length. Thin glass plates can be

bent to radii greater than 18 inches and simple

field flatteners yield satisfactory results in

cameras with focal ratios up to about d/f=l.

However, to exploit fully the still higher

speeds in the shorter focal lengths it will be

necessary to develop either field flatteners that

are effective at the higher focal ratios or to find

a flexible emulsion base that is stable against

changes in temperature and humidity.

Another urgently needed development is an

optical system that will provide a wide field of

1

 Director of Mount Wilson and Paloraar Observatories.

critical definition for telescopes of very large

aperture. Schmidt telescopes have been successfully made with apertures up to 48 inches

with a focal length of 120 inches. However, in

order to avoid chromatic aberrations the focal

length must increase as the 3/2 power of the

aperture. Thus the minimum focus of a 200-

inch Schmidt is 1,000 inches. The length of

the tube would have to be twice this, or 2,000

inches; i. e., three times that of the present

Hale telescope. This would make the cost of

the telescope tube and the dome to house it

prohibitive.

Unfortunately the field of good definition of

a high-speed paraboloid is very small; less than

}{ inch in diameter for the Hale telescope in

which d/f= 1/3.3. Corrector lenses designed

by Dr. F. E. Ross and placed in front of the

plate have successfully increased this area of

good definition to a diameter of 3 or 4 inches.

This still falls far short of the 14 X 14-inch area of

critical definition attained with the 48-inch

Schmidt, which operates at the even higher

speed of d/f= 1/2.5. Because of this small

field, the use of large telescopes such as the

200-inch is impractical for most survey purposes.

Another important need at the present time

is an objective prism or similar device that may

be used for the wholesale classification of spectra

of stars of magnitude fainter than can now be

reached. Thus, because of the deterioration of

the definition of the spectra by "seeing," it is

not feasible to use an objective prism on telescopes of appreciably more than 10-foot focal

length. Furthermore, the limit on the exposure

time set by the sky background is essentially

the same with an objective prism as for direct

photography. If the usual dispersion of 200-

300 A/mm is used and the spectra are widened

to 0.2 or 0.3 mm, the starlight is spread over an

area 1,500 to 4,000 times that of a star image,
SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS

with the result that the limiting magnitude for

the spectra is 8 or 9 magnitudes below the limit

for direct photography. For a 10-foot focal

length the limiting magnitude for direct photography isMph=21 and therefore for spectroscopy Mph=12 or 13. This falls far short of the

magnitudes that may be reached with slit

spectra. Thus with the 200-inch telescope a

1-night exposure enables one to reach the 15th

magnitude at 38 A/mm and the 18th magnitude

at 170 A/mm. At the present time there is no

equipment available for picking up unusual

faint objects for detailed study at these higher

dispersions with a slit spectrograph. The anticipated development of the image tube, discussed

below, will render this problem even more serious

by pushing the limit of the slit spectrograph

from 2 to 5 magnitudes still fainter.

One suggested solution of this problem is the

use, as a multislit, of a high-contrast positive

print of a given star field taken with the same

telescope. With a wide-angle telescope and a

spectrograph having wide-angle collimator and

camera optics, a large number of spectra could

be obtained simultaneously without interference

from the sky background. Unfortunately, if

Schmidt optics are used throughout, the field

flatteners necessary to bring the curved fields of

the telescope and collimator into coincidence

destroy the off-axis collimation. Some other

wide-angle system must be devised for the

purpose.

The efficiency of slit spectrographs would be

substantially increased if larger gratings could

be made or if gratings could be ruled which provide higher angular dispersion and still retain a

high concentration of light in one order. In

order to avoid interference between collimator

and camera optics, this grating should preferably be of the transmission type. The so-called

"venetian-blind type" described in "Astronomical Spectrographs: Past, Present, and Future,"

in ' 'Vistas in Astronomy,'' is one possible solution

if the necessary ruling techniques can be

developed.

Another problem that needs investigation is

that of seeing. Presumably very little can be

done to reduce turbulence in the upper atmosphere. However, in solar instruments much of

the disturbance is caused by the heating of the

optical parts and their surroundings by the sunlight which is incident on them. Undoubtedly

the local air turbulence due to this could be substantially reduced by proper light shields and

by artificial cooling of parts that cannot be

shielded. Useful studies could also be made to

determine the most effective method for thermostating an optical instrument such as a spectrograph without setting up internal air currents

and turbulence that often cause serious deterioration of the image.

Every large telescope is essentially a collector

of light which is focused on some type of

receiver. The properties of this receiver play

a very large role in the proper design of the

telescope. Thus up until the present century

practically all astronomical observations were

made visually; i. e., the human eye was the

receiver. The range of wavelengths to which

the eye has appreciable sensitivity is very small

(1,000-1,500 A) and the field of view of critical

definition is also very limited. Furthermore,

the lens of the eye is a part of the optical

system and sets an upper limit to the size of

the cone of light that can be concentrated on

the sensitive surface. For this type of receiver

the simple two-lens refractor operating at focal

ratios of d/j'=1/15 to 1/20 was ideal. As a

result nearly all large telescopes constructed

in the last half of the 19th century were of this

type.

By 1900 photographic procedures were rapidly

replacing visual observations. The silver-halide

grains, however, are sensitive in the ultraviolet

to beyond the limit of transmission of the

atmosphere while the more recent development

of dyes has produced plates that are sensitive

well out into the infrared range. Also, in order

that an image be recorded on a photographic

plate it is necessary that a certain minimum

number of quanta fall on a unit area of the

plate. This limit is such that in order to

photograph faint galaxies and nebulosities in

a feasible exposure time it is necessary to use a

light collector with a focal ratio of at least

djf=l/5. This fact resulted in a sudden shift

early in the present century to the use of the

reflector, which provides the complete achromatism and greater speed necessary for the

efficient use of the new type of receiver.

Not only were the properties of the photographic plate major factors in causing the shift
NEW HORIZONS IN ASTRONOMY

from refractors to reflectors, but they have

played an important role in fixing many of the

design criteria for both the reflector itself and

the auxiliary equipment, such as spectrographs,

to be used with it. One of these properties is

the linear resolving power of the plate. This

resolving power sets the minimum focal length

of the main telescope or of the spectrograph

camera if a given angular resolving power or

wavelength resolution, AX, is to be achieved.

Furthermore, in all of these instruments this

plate resolving power is the primary factor

considered in setting up optical design and

construction tolerances. Another property of

the plate, namely, that it can record a very

wide field at one time, led to the development of

various wide-field optics culminating in the

Schmidt camera. Similarly the necessity of

concentrating a certain minimum of light per

unit area in order to obtain an image has been

the chief cause of the design and construction

of optical systems of extreme speed («?//= 1 or

greater) for spectrograph cameras.

It is now becoming apparent that another

change in the receiver for astronomical observations is in the offing. Thus photoelectric

surfaces have been developed, and are in use

for television and other purposes, that have a

quantum efficiency from 10 to 100 times as

high as that of the best photographic plates.

The electrical circuits used in some of these

receivers permit the reduction or elimination

of a uniform background of the type caused by

scattered starlight and the permanent aurora,

which sets the present limit in direct photography of the sky. Indeed it now appears that

in the foreseeable future the limit of detectability of a faint star image will approach the theoretical limit. This limit is set by the necessity

of collecting a larger number of quanta from

the star than the probable random fluctuation

in the number of quanta received from the sky

background in the area of the sky covered by

the object.

These developments are still in such an early

stage that it is too soon to predict their possible

effect on telescope design. Thus the linear

resolving power will depend on the final type

of receiver adopted. It is also quite probable

that the new receiver will not require, like the

photographic plate, a minimum concentration

of quanta per unit area for a satisfactory record.

One criterion for this will depend on the possibility of eliminating spurious thermal electrons

by refrigeration or other means. In any case

the much greater efficiency of the photoelectric process will permit the record of a

given object to be obtained in very much less

time than that required by the photographic

process. The shift in these two properties will

probably eliminate the necessity for optics of

extreme focal ratio such as the types recently

developed for both direct and spectroscopic

observations. Regardless of the properties

of the photoelectric receiver as finally developed,

substantial changes in telescope and spectrograph design may be expected, if for no other

reason than the elimination of many of the

present restrictions set by the properties of the

photographic plate.

Solar Instrumentation

By J. W. Evansx and R. B. Dunn

The technical problems of solar observation

are decidedly different from those in any other

field of astronomy because of the abundance

of light and the large apparent diameter of the

sun. The latitude for the exercise of ingenuity

in securing more and more refined information

is practically unlimited, and affords a kind of

satisfaction rarely found in any other field.

The next 25 years promise to be the most

exciting and the most productive since the formulation of the laws of spectrum analysis by

Kirchhoff. Since 1945, the conception and

construction of new instruments and techniques

have accelerated markedly, as have the theoretical advances which designate the critical

observations we need to determine the physical

character of the sun. We have barely begun

to get acquainted with the new tools and to

exploit their possibilities, but the results are

already impressive.

Most of the new developments are projections

of older devices and methods, with refinements

that greatly increase their power. Two of

them, however, introduce entirely new methods

that are sure to be exceedingly fruitful: the

observation of the far ultraviolet spectrum of

the sun from rockets, and the recording of solar

radio emission.

The rocket observations are a fine example

of cooperative effort in several fields to solve a

very difficult technical problem. The recording instrument is either a spectrograph designed

to work in the 900- to 3,000-angstrom range or

an X-ray ionization chamber for the 1- to

5-angstrom region. The recorder must be

mounted in the rocket with a device which keeps

it pointed at the sun with an accuracy of a few

minutes of arc, in spite of lively gyrations in

its supporting platform. Finally, the rocket

itself must be capable of carrying its load to

an altitude of 100 km or more in order to rise

1

 Sacramento Peak Observatory, Sunspot, X. Mex.

above the regions of the earth's atmosphere

that absorb the solar radiations in these wavelength regions and prevent their observation

from the ground.

This brief description can give no concept of

the difficulties encountered in the work, nor

of the years of patience and imagination that

were required to overcome them; but the results so far have more than justified the effort.

The expected spectroscopic features duly appeared, but some very unexpected emission

lines were also recorded. The expected lines

strengthen our confidence in the broad features

of existing solar theory, and the new lines

present new conditions which any valid theory

must satisfy. The rocket observations have

unveiled a hitherto inaccessible part of the

solar spectrum, but the quick glance we have

been allowed sharpens the appetite for more

details, more quantitative measurements, and

a further extension to still shorter waves.

Rockets can still do a great deal, but there

can now be little doubt that a man-made satellite eventually will be our platform for solar

research above the absorption of the terrestrial

atmosphere. New difficulties will appear and

be overcome, and new data on the solar

spectrum will further delineate the structure

of the sun. Such advances, inevitably, will

raise new questions in place of the old ones

solved, but they will concern questions of details which we cannot ask at present because

we do not know enough.

It was discovered during World War II

that the sun was the source of occasional

radio noise strong enough to interfere with

radar operations. This discovery has led to

a new branch of solar astronomy which we have

just begun to explore in its general outlines.

Instead of the 2-to-l range in wavelength of

the visible spectrum, the observable radio

emission from the sun covers a range of at
SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS VOL.1

least 1,000 to^l. This provides a tool for

analysis in depth. Shortwaves in the centimeter range originate near the visible surface

of the sun. As we go to longer waves, the

point of origin rises through the corona to

heights of some hundreds of thousands of

kilometers for the longest waves of 10 meters

or so. Like the light spectrum, the nature

of the radio spectrum depends on physical

conditions at its point of origin and in the

higher layers traversed by the radiation in

leaving the sun. The interpretation of much

that has been observed is not yet clear, but

there can be no doubt that solar radio astronomy

will occupy an extremely important place in

future research.

The design and construction of new solar

instruments at observatories all over the world

continues at an ever-increasing pace. Some

of them, like the Babcocks' magnetograph,

are the response to a particular need. Others

are new versions of tried-and-true devices,

like spectrographs, for which there are sure to

be many uses not necessarily specified in advance. In a field where we have still so much

to learn, we should plan our new equipment

with care in order to avoid unprofitable duplication.

In considering solar instrumentation, both

present and future, we should always keep ia

mind our real object: to find out something

about the sun. Specifically, we want the statistics, photometry, and physics of the solar

atmosphere, and a good theoretical model

that explains all the observations. We must

approach the problem from every possible

direction. Routine daily survey programs,

detailed studies of active regions, critical

observations designated by theoretical models,

new types of observations—all should be used

to solve the problem.

Each new instrument acquired should fill a

specific need. A careful investigation is often

required to determine whether the planned

observation is already being carried out elsewhere. Especially in planning a survey program, the observer should ask himself: Should

I make an effort to take a few prominence

movies, when other observatories have collected

such gross quantities of film that they defy

any sort of reduction ? Will my daily survey of

sunspots contribute significantly to the daily

coverage obtained elsewhere? Can the solar

activity be adequately described from present

solar surveys? In other words, the observer

should have a specific problem in mind when

building the instrument and when taking the

observations.

In instrumental development there is room

for every type of genius. To get an ingenious

idea and show that it works is not enough.

The inventor, or someone else, must still

develop the instrument to the point where any

capable observer can use it routinely to obtain

reliable data. A typical example of instrument development is the coronagraph. Lyot

showed that the coronagraph could be built,

and he proceeded to build one. He used it

with the utmost skill to solve several important

coronal problems, and, incidentally, demonstrated its capabilities. Since then a great

deal of work has resulted in developing the

instrument, as a coronagraph, and in making

it more convenient to operate. It is hoped

that due to additional scale, resolution, photometry, and systematic use, Lyot's ideas will

tell us even more about the sun than did his

first observations. In some cases the conveniences themselves make the difference between taking or not taking the critical observation.

The development of the completed instrument is very often the result of a team effort.

We need new gadgets and new ideas, but in

the solar field we also need people who take

existing ideas and instruments and add such

vital factors as photometric standards, direct

photoelectric photometry, differential photometry, seeing cancellation devices, and conveniences that speed the operation.

As elsewhere, solar instrumentation is getting

so complex that we need specialists in the optical, electronic, and mechanical aspects and in

data reduction, who will cooperate with the solar

physicist to produce the final result. It is no

longer feasible for a person to be skilled in all

fields necessary to the solution of a problem.

He must work with a group, and the group

members must have mutual confidence in each

other. A great deal of intercommunication is

essential. Finally, the group must be competently directed and coordinated.
NEW HORIZONS IN ASTRONOMY

The value of some very simple solar observational device in an educational institution

should never be underestimated. An example

is a prominence birefringent filter. The direct

contribution to research on the sun may be small

or none, but the stimuli provided by such facilities may be decisive in attracting capable students to solar research, which greatly needs

skilled workers.

Until a few years ago solar research was mainly concerned with broad average characteristics.

One studied the spectrum without regard for

small variations from point to point on the sun.

The chromosphere was treated as a homogeneous atmosphere of uniform temperature. More

recently the importance of the local differences

has been recognized and the tendency has been

toward a finer grained analysis, in which many

small areas of the solar surface must be studied

separately. The result is an enormous increase

in the volume of information required, since we

now have to consider many small areas instead

of a single large one. Gathering data, reducing

them to useful numerical quantities, and analyzing the greatly increased volume of such numerical quantities are already acute problems,

and will become more serious as time goes on.

In gathering data, the unaccustomed pinch

of the basic limitation of information contained

in the available light from the sun is already

being felt. A telescope receives light from a

given area on the sun in the form of discrete

quanta at an average rate of n per second, in a

purely random fashion. This means that the

number of quanta collected by the telescope in

a given second differs from n in a random fashion, with a mean percentage error of -j=-

Hence a 1-second observation can at best determine the brightness of the area with this

percentage accuracy, and the larger n is, the

more accurate the observation. When we try

to observe smaller and smaller areas of the sun

we are, in effect, reducing n in proportion, and

the fundamental errors of observation increase.

Unfortunately we cannot make up for this by

simply observing for a longer time. Seeing is

never good enough to yield an undistorted and

unsmeared image of the sun for more than a

fraction of a second, and the longer our observation the more our small area becomes confused

with light from neighboring areas. The idea of

a scarcity of sunlight is perhaps a novel one, but

may not seem so surprising when we consider

that we would like right now to take the light

from an area of 0.01 square millimeter in a

300-mm solar image, spread it out into a visible

spectrum nearly a hundred yards long, and

then determine the intensity in an area of 0.01

square mm of this spectrum. The one thing we

can do is to increase efficiency in the use of the

quanta we receive. The photographic plate

utilizes only one in a hundred of the incident

quanta. Modern photoelectric surfaces are

about 10 times as efficient, and we have already

begun to take advantage of this improvement.

The problem of the future is to find still more

efficient photoelectric surfaces, and to develop

accurate and convenient methods for recording

continuously the photoelectrons emitted from

each point of the surface.

Once the observational data are secured, the

problem of reducing them to numerical quantities must be faced. This has always been a

laborious process, and the vastly greater volume

of data now required demands a radically new

method of automatic reduction. The magnificent spectra coming from the new vacuum

spectrograph at the McMath-Hulbert Observatory emphasize the problem decisively, and I

take them as a clear-cut example. The

McMath-Hulbert observers find that both the

Fyaunhofer lines and the solar continuum

show marked variations in the direction

perpendicular to the dispersion. These variations are due to spectroscopic differences in

the small regions imaged on the slit of the

spectrograph. They consist in wiggles and

changes of width in the lines, and changes in

the brightness of the continuum. Thus the

spectrograms have variations which must be

measured in two directions instead of one, along

the spectrum and across the spectrum. Now

suppose we want to exhaust the information

contained on such a plate over a wavelength

range of only 5 angstroms. Let us say that the

slit was 1 cm long, that the resolving power of

the telescope was 0.2 mm along the slit, and

the resolving power of the spectrograph was

0.01 A. Complete information, therefore, calls

for the determination of the light intensity at

500 points along the spectrum for each of 50
s

SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS

points along the slit. In other words, we have

to make 25,000 intensity determinations for a

complete analysis, instead of the 500 that would

suffice if there were no appreciable variations

along the slit. In this instance the new

dimension has multiplied the labor of analysis

by a factor of 50. The information is there in

the plate, but the labor of extracting it has

become prohibitive.

The case of the wiggly lines is only one example of the trend toward the study of finer

details on the solar disk, in the chromosphere,

and in the corona. The result is an unprecedented mass of information in the coded form

of photographic density, magnetic variations

in a tape, or some other analog record. To be

of any use, these records must be decoded and

presented in numerical form. Once this has

been accomplished, we are confronted with

endless pages of numbers, all ready for interpretation by the unhappy theoretician. Classical

methods of data handling will be quite inadequate to utilize more than a tiny fraction of

the information contained in the observations,

and future progress depends very heavily on

finding the means for breaking this bottleneck.

These problems of data reduction and analysis

on a large scale are by no means limited to

solar astronomy. Industrial and military needs

have stimulated the development of information

theory, the science of extracting the maximum

possible information from a minimum of raw

data in the least time. Devices for automatically performing various standard operations in

data reduction, and computers which handle

the numerical results with unimaginable speed,

have been devised. The theoretical and physical machinery for the solution of the astronomical problem are now at hand, and the computers at least are already vigorously at work for

astronomers in several fields. But as yet the

automatic reduction of astronomical data to

the point where they are ready to be fed into

a computer has hardly been touched. The

problem for the immediate future, therefore, is

to develop the devices for bridging the gap

between the output of a telescope and the input

to a train of automatic reducing equipment

already available. The output of such a train

is numerical information, ready either for the

computer or for direct analysis by the astrophysicist if the quantity is manageable. The

importance of this problem should be emphasized. Unless it can be solved the effectiveness of many of the instrumental developments for solar research will be limited to a

fraction of their real potential.
Astronomical Seeing and Scintillation1

By Geoffrey Keller

Anyone who has made observations with large

telescopes is only too painfully aware that the

most serious limitations to their performance are

usually connected in one way or another with

problems of astronomical seeing. Images of

stars are much larger than one would expect

them to be were their size determined only by

the aberrations of the telescope itself. Details

of extended objects, such as the moon and the

sun, are badly smeared when the seeing is poor.

As a case in point, Baum (1955) gives the average image diameter seen with the 200-inch telescope at about 2.5 seconds of arc, whereas in the

absence of bad seeing conditions the perfectly

adjusted telescope should give images of diameter of the order of 0.04 seconds of arc. The

total amount of light received from a bright star

by a telescope of moderate size also fluctuates

in an irregular fashion, and in a manner which

can easily be shown to be due not to intrinsic

fluctuations in the star but to the fact that the

starlight must pass through regions of the

earth's atmosphere wherein the density varies

irregularly. Protheroe (1955a) has shown that

with a telescope of 12% inches aperture the

fluctuations in brightness of a bright star will

average around 10 percent of the mean, and will

have frequencies ranging up to a thousand cycles

per second. In pioneering investigations in this

field Mikesell, Hoag, and Hall (1951) obtained

similar results.

Similar effects are encountered in radio astronomy, and here also the responsible agency

is a region of nonuniform index of refraction in

the earth's atmosphere. Whereas in the case of

optical scintillation the principal effects apparently are caused in the troposphere and lower

stratosphere, the majority of the effects on radio

1

 "Scintillation" is used here to mean only the changes with time of

the brightness of the telescopic image of a star or other source. "Seeing"

is construed to mean the degree of change in the position, shape, or size

of a stellar linage or of a detail of an extended object.

1

 Perkins Observatory, Delaware, Ohio.

waves seem to be caused by nonuniform electron

densities in the ionosphere. Ryle and Hewish

(1950) have observed fluctuations in the intensity of radio sources at meter wave lengths of

the order of 10 percent and fluctuations in the

apparent positions of the radio objects of the

order of a minute or so of arc. Such variations

give rise to serious observational problems, and,

though some observational uncertainties due to

scintillation and bad seeing may be overcome by

making repeated measurements, the loss in efficiency is very great.

The most extensive efforts that have been

made to date to minimize the effects of the

earth's atmosphere on astronomical observations have consisted of more or less empirical

comparisons of various possible sites for establishing new observatories. When a very large

investment of this sort is to be made, it becomes

economically feasible to spend a considerable

sum on extended tests. Such tests were made

prior to the construction of the Hale telescope

at Mount Palomar, and currently a new series

of tests is being made in connection with a

search for a site for a proposed American Observatory. Such surveys generally use telescopes and record the fluctuations in size and

brightness of a star such as Polaris (chosen

frequently because of its relatively fixed position

in the sky).

Another potential means of reducing seeing

effects is through the use of instruments mounted on high-altitude balloons, rockets, and eventually extraterrestrial satellites. These methods

might eliminate our seeing problems, but before

adopting them we should know, for example,

how high the instrument platform must be

in order to be safely above the troublesome

layers of the atmosphere. In the case of a

radio telescope, the answer is likely to be of

the order of hundreds of miles. There are

obvious difficulties with such an approach

9
10 SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS VOL. 1

when we attempt to apply it to large telescopes,

so that it will probably be many years or even

centuries before the bulk of astronomical

observations can be made at extra-atmospheric

stations.

Having picked a site where the disturbances

caused by intrinsic seeing and scintillation are

at a minimum, the astronomer is next interested

in knowing how best to design his telescope

and observatory so that thermal disturbances

in the air caused by their structure may be

minimized. A number of theoretical and observational studies of the behavior of scintillation have been made at the Perkins Observatory that seem to point more and more to the

conclusion that the local thermal irregularities

of the air may be the predominant cause of

bad image distortion (as distinguished from

bad scintillation). The work of the astronomers

at Paris and Haute-Provence seems to have

led them to similar conclusions, so that the

new 193-cm telescope for the observatory at

Haute-Provence is being very carefully designed to minimize the temperature irregularities in the air in the tube and just outside

the dome (see Couder, 1953). Should these

efforts be successful, a major improvement in

the performance of large telescopes may result.

The importance of minimizing irregular

refraction in the gas inside of auxiliary astronomical equipment has been well demonstrated

by the experience of observers with the new

52-foot vacuum spectrograph at the McMathHulbert Observatory (McMath, 1955). The

steadiness and sharpness of the spectral lines

were greatly increased as the spectrograph was

evacuated. Doubtless there are many other

pieces of auxiliary equipment whose performances could be greatly improved by a similar

procedure.

Although the writer knows of no successful

instruments, so far devised, to compensate for

(rather than eliminate) the effects of seeing,

some have been envisioned and might well be

investigated (Babcock, 1953). Any device which

attempts to compensate for seeing must depend, first, on knowledge of what sort of optical

irregularities exist in the air; this requires study

of the distortion in the wave fronts in starlight

which has just passed through a layer of irregularities. The device must then make, presumably by electro-optical methods, compensating

adjustments in tie optical system in the telescope. Thus, if the air in a portion of the light

path is unusually warm, the corresponding index

of refraction will be abnormally low and the

velocity of light abnormally high. As a result,

a wave front tends to advance farther than it

should as it passes through this region. The

compensating device attempts, in effect, temporarily to move the surface of the primary mirror

farther back in the area where this portion of the

wave front is reflected, so that after the reflection of the wave front the irregularities have

been ironed out. Methods of this type have

certain drawbacks. For example, we must observe objects that are sufficiently bright to give

enough photons per second to activate the

device which controls the compensator; yet,

enough photons must be left over to make

possible the primary astronomical observation.

In the preceding paragraphs we have looked

at the problems of astronomical seeing from the

point of view of the typical astronomer, to

whom seeing is an unmitigated evil that must

be avoided or removed by whatever means come

to hand. To the atmospheric physicist and

meteorologist, the situation may be quite the

reverse. The fluctuations in the light of a

bright star, known as scintillation, are known to

be caused by the motion across the telescope

objective of a complex pattern of lights and

darks, known as the shadow pattern or shadow

bands. This shadow pattern is caused by the

passage of starlight through atmospheric irregularities which must occur at a considerable

height. These irregularities diffract the light

and cause interference and reinforcement of the

wave front at various points along the ground.

A fairly complete theory of the relation between

the atmospheric irregularities and the pattern

of lights and darks in the shadow pattern has

been developed by van Isacker (1953). With a

reliable means of studying the structure of the

shadow pattern, much may be said about the

pattern of the density variations in the air itself.

A suitable method for observing the shadow

pattern has been developed by the author

(1955) and by Protheroe (1955b). There seems

to be considerable promise that by the use and

extension of this method we might be able to

obtain considerable information about the
NEW HORIZONS IN ASTRONOMT 11

nature of clear air turbulence, its motions, its

height, and how these quantities vary with

meteorological conditions. At the same time,

more empirical studies of the relationship

between scintillation and wind velocity are

being made (Gifford and Mikesell, 1953;

Gifford, 1955; Protheroe, 1955a; and Mikesell,

1955).

Studies of the scintillation of radio objects

are also being made, and much information

about the structure of the ionosphere is being

obtained (Pawsey, 1955).

To summarize, the study of astronomical

seeing and scintillation is of vital interest to

the astronomer in the same sense that studies

of optics and radio technology are of importance

to him. There is much to be learned about the

subject, and there is considerable promise that

with more understanding will come ways and

means of greatly improving the performance of

telescopes. Although seeing may be a nuisance

to the astronomer, it also is a tool for making

both fundamental and routine studies of the

earth's atmosphere.

A few observatories in this country and

abroad are making studies of both kinds.

Much exciting work remains to be done and it

is very much to be hoped that additional

astronomers can be persuaded to take an

interest in the subject.

References

BABCOCK, H. W.

1953. Publ. Astron. Soc. Pacific, vol. 65, p. 229.

BAUM, WILLIAM A.

1955. Sky and Telescope, vol. 14, pp. 264, 330.

COUDER, ANDRE.

1953. Publ. TObe. Haute-Provence, vol. 2, No. 53.

GIFFORD, FRANK, JR.

1955. Bull. Amer. Meteorol. Soc., vol. 36, p. 35.

GIFFORD, F., AND MIKESELL, A. H.

1953. Weather, vol. 8, p. 195.

KELLER, G.

1955. Journ. Opt. Soc. America, vol. 45, p. 845.

Also, Contr. Perkins Obs., ser. 2, No. 5.

MCMATH, R.

1955. Sky and Telescope, vol. 14, p. 372.

MIKESELL, A. H.

1955. Publ. U. S. Naval Obs., ser. 2, vol. 17, pt. 4.

MIKESELL, A. H., HOAG, A. A., AND HALL, J. S.

1951. Journ. Opt. Soc. America, vol. 41, p. 689.

PAWSEY, J. L.

1955. Radio star scintillation due to ionosphere

focusing. Physics of the ionosphere.

The Physical Society, London.

PROTHEROE, W. M.

1955a. Contr. Perkins Obs., ser. 2, No. 4.

1955b. Journ. Opt. Soc. America, vol. 45, p. 851.

Also, Contr. Perkins Obs., ser. 2, No. 6.

RTLE, M., AND HBWISH, A.

1950. Monthly Notices Roy. Astron. Soc. London, vol. 110, p. 381.

VAN ISACKER, J.

1953. Publ. Inst. Roy. Me'teorol. Belgique, ser. B,

No. 8.

Spectroscopy

By Charlotte E. Moore

The outlook for future research in spectroscopy is highly promising. The rocket

spectroscopy of the solar spectrum, for example,

extends our horizon to the shortwave region,

and thus provides a far more complete picture

of solar spectroscopy than ever before. The

O vi solar emission lines and other observations

of equally far-reaching significance demonstrate

the astrophysical importance of further study of

high-ionization spectra in the laboratory (Johnson, Malitson, et al., 1955). Ultraviolet standards of wavelength must be established in the

shortwave region, and more term values must

be worked out, particularly in spectra of the

more abundant elements. Fe iv is an outstanding example of a neglected spectrum urgently

needed by the astrophysicist. The coronal

spectrum, the nebular spectra, and now the

spectra of distant galaxies, whose lines are

shifted further and further to longer waves, all

show that further study of ultraviolet spectra is

extremely important.

In the longwave region, beyond the photographic range, the development of suitable heat

detectors has extended the astrophysical vista,

as well as that of the laboratory, to the far

infrared, and microwave spectroscopy has

extended it still further. With the astrophysical range from Lyman alpha (1215 A) to the

21-cm wavelength of hydrogen, our studies can

no longer be limited to selected regions of the

"visible" spectrum.

Among high-dispersion stellar spectra the

sunspot spectrum is one of the most challenging.

Hundreds of spot lines have not yet been

measured; many of them are doubtless of

molecular origin. Furthermore, no homogeneous set of sunspot spectrograms exists.

During the coming sunspot maximum, highdispersion sunspot spectrograms should be

taken with the same spot, over the whole

1

 National Bureau of Standards, Washington, D. C.

spectral range. This is a taxing but vital

program.

Our knowledge of stellar spectra is far from

complete. Any study of abundances requires

correct identifications of the lines, together with

a knowledge of the excitation and ionization

energies of the separate atoms and ions.

P. W. Merrill has recently reported that in the

Se star, R And, there are some 275 unidentified

absorption lines between 4000 A and 6900 A.

J. Greenstein has observed several hundred

unidentified lines in v Sgr. (see Merrill, 1956).

In the solar spectrum about 30 percent of the

lines are still of unexplained origin. These and

other similar examples demonstrate how much

work is yet to be done in spectroscopy. This

need is reflected in the persistent demand,

among astronomers, for a new, large, revised

multiplet table that extends from the far

ultraviolet to the extreme infrared.

These demands impose upon the laboratory

spectroscopist the task of observing and

analyzing more atomic spectra, and also of

selecting spectra to be studied that anticipate

future astrophjsical needs. The use of modern

sources, such as electrodeless discharges and

infrared detectors, is one of the most promising

approaches to the subject. Earlier descriptions

of atomic spectra are not only lacking in faint

lines but are incomplete, also, because of masking by oxide bands. The electrodeless discharge

has the advantage in atomic work that it can

be controlled in such a way as to suppress

molecular lines. More than 1,200 Fe i lines

have been identified in the solar spectrum by

prediction from the known atomic energy levels.

A suitable laboratory source of excitation would

doubtless excite these and many more faint

lines. Kiess has already observed a number

of the predicted Fe i lines in selected regions.

EdleVs (1953) observations of CA n in the

infrared have resulted in the identification of

13
14 SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS

four new multiplets among solar lines. For

many familiar spectra there exists no complete

or homogeneous set of observations. Modern

equipment and techniques could be used to

good advantage in improving this situation.

The spectra of the rare earths are the most

conspicuous group that needs study. Within

the next 3 years it is planned to enlist the

services of all available experts in attacking

this difficult problem. Volume 4 of "Atomic

Energy Levels" will be devoted entirely to

rare-earth spectra. This program should provide the astrophysicist with a wealth of new

data. In order to disentangle the intricacies

of these spectra, the assistance of theoretical

investigators to predict the positions of the

terms of various configurations is needed.

This problem leads to use of digital computers

in solving large matrices. Here is a useful field

of activity in theoretical problems that parallels

the laboratory and astrophysical research.

The comments just given apply chiefly to

atomic spectra, but a wide vista appears, also,

in molecular spectroscopy. For example, the

use of controlled flames as a source makes it

possible to bring out certain band spectra and

suppress others. Precise laboratory measurements and analyses, particularly for diatomic

molecules of the more abundant elements, are

sadly needed. The writer is continually receiving requests for a critical compendium of

molecular data paralleling the volumes on

"Atomic Energy Levels." In solar and spot

spectra alone this subject is far from completed

and promises interesting results in future

astrophysical research.

Spectroscopy embodies, also, the study of

line intensities, equivalent widths, and curves

of growth—all rewarding subjects for future

research, particularly with high-dispersion stellar spectra. Just as in the case of spectrum

analysis, measured laboratory intensities are

required to increase our knowledge of abundances of elements. This phase of the work

must be carried on for years to come.

The door is open and the opportunity is rich

for the student interested in spectra. The

more unusual examples of astrophysical identifications—nebular lines, coronal lines, solar

emission lines of high ionization, and technetium

lines in the S-stars—all challenge the astrophysicist and in themselves more than justify

active financial support of spec troscopic research

in the laboratory, in theoretical fields, and in

the observatory.

References

EDL£N, B.

1953. Accad. Naz. Lincei Conv. Sci. Fis. Mat.

Nat., vol. 11, p. 56.

JOHNSON, F. S.; MALITSON, H. H.; PUBCELL, J. D.;

AND TOTJSET, R.

1955. Astron. Journ., vol. 60, p. 165.

MERRILL, P. W.

1956. Carnegie Inst. Washington, Publ. 610.
Radial Velocities

By R. M. Petrie1

The measurement of the line-of-sight component of stellar motion ranks among the early

achievements of stellar spectroscopy. The

principle was announced by Doppler in the

middle of the 19th century and was extended

and elaborated by a number of astronomers and

physicists soon after. The first applications

were made by visual methods and are now of

historical interest only. The subject did not

attain full stature until the introduction of

photographic methods during the closing years

of the century, but it then became at once a

highly successful and valuable technique of

astronomy. Although radial-velocity work is

thus about 60 years old, it is still of prime importance and promises to remain an indispensable tool in many branches of astronomy.

The pursuit of this work is not successful without a considerable expenditure of patient

effort, but the student may look to it to bring

substantial rewards for his labors. He may,

with profit, reflect upon the prophetic words of

Huggins: "It would scarcely be possible, without the appearance of great exaggeration, to

attempt to sketch out even in broad outline the

many glorious achievements which doubtless lie

before this method of research . . . . "

In this article attention will be directed to

present activities and some probable future

needs. The development of our subject from

the beginning, outlined above, is described in

the first part of Campbell's "Stellar Motions,"

published in 1913.

Spectrographs

After the introduction of photography the main

question to be answered was how accurately the

spectrograph could record the positions of

stellar absorption lines with reference to some

terrestrial comparison source. The Doppler

1

 Director, Dominion Astrophysical Observatory, Victoria, British

Columbia, Canada.

shift is always small and great care is necessary

to avoid spurious instrumental shifts such as,

for example, those introduced by mechanical

flexure, temperature variations, and faulty

optical components. The development of reliable instruments such as those at Potsdam,

Mount Hamilton, or Victoria may be studied

in an article by Eberhard (1933). The main

problems of design have now been solved, and

there exists a considerable number of instruments capable of yielding accurate radial

velocities.

Some problems of instrumentation remain as

a challenge to the student. The following are

examples of these.

Gratings.—The development of the "blazed"

grating of high efficiency opens a field for application to radial-velocity work in the visual

region. Spectral types of F 5 and later may be

studied effectively in this way, which offers two

important advantages. First, the spectra are

simpler in the visual region, so that one can

minimize the troublesome problem of using

"blended" lines; and, second, many stars are

brighter in the visual than in the usual photographic region. A successful grating spectrograph has been constructed and used by Merrill

(1931), and of course the powerful coud6

spectrographs of the largest telescopes employ

gratings as the dispersive units. There is still

room, however, for experiments in the most

effective use of gratings in radial-velocity work

with moderate dispersion.

Long-focus instruments.—The great power of

modern coude spectrographs lends itself well to

radial-velocity investigations requiring the highest precision. It is not yet known with certainty whether the long-focus collimator and

wide-slit spectrograph fed by a combination of

three to five mirrors realizes the potential accuracy of the very high dispersion. We want

to know, on the one hand, whether the wide

15
16 SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS

slit (relative to the stellar image) has limited

the accuracy because of the introduction of

guiding errors; on the other hand, we want to

know whether the multimirror system has preserved full "identity of source" between stellar

and comparison spectra. Experimental work

to answer these questions needs to be undertaken as a preliminary to radial-velocity work

of the highest precision. As an example of the

inherent possibilities, one may refer to Adams'

(1941) spectroscopic determination of the solar

parallax.

Comparison spectra.—Present practice generally is to use an arc, or spark, in air, with the

most convenient metals being iron or titanium.

These sources are fairly satisfactory but are not

entirely constant unless a number of precautions

are taken. It is possible that a better source

can be found in some modification of a discharge

tube such as the hollow cathode type recently

described (Edlen, 1955).

The greatest obstacle to the study of stellar

motions through radial velocities is the slow

rate at which spectrograms and measures accumulate. The slit spectrograph, while accurate, is very wasteful of starlight and ordinarily

utilizes only a few percent of the incident light.

As a result we are, for example, almost totally

ignorant of the radial velocities of stars of

apparent magnitude 10 and fainter. Fifty

years of work have yielded velocities for some

15,000 stars, but this sample is much too small

for galactic studies, and current progress is slow.

A major effort and a new attack are obviously

required if the study of stellar dynamics is not

to be hampered by lack of observations.

Fortunately there is now some hope of improving matters through two different current

developments.

Slitless spectrographs.—The idea of removing

the spectrograph slit and collimator and thus

photographing with one exposure the spectra of

of a field of stars is an old one. An account of

the early experiments with objective prisms is

given by Millman (1930).

Until recently the objective-prism work fell

short of its goal because of the presence of field

distortions and instrumental errors exceeding

many times the Doppler shift. However a

series of important experiments by Fehrenbach

(1947, 1948) have resulted in the development

of an objective prism free from distortion, and

studies by Treanor (1948) and Schalen (1954)

have verified and extended this result. A

wholesale accumulation of radial velocities of

faint stars is now in prospect; continued experimentation and application are sure to bring

rich rewards. Already the "mean errors" have

been reduced to about ±6 km/sec, and this is

sufficient for application to many problems in

galactic motions and structure.

Image convertors.—Another promising development is the image convertor, which is

capable of high efficiency and may ultimately

replace the use of photographic plates in recording stellar spectra (Argyle, 1955). Considerable experimentation is now going on to

improve this technique and to apply it to stellar

spectroscopy. The subject is well worth the

attention of the student who wishes to make

a contribution to radial velocity work.

Measurement

Doppler shifts are, generally speaking, very

small and their successful measurement calls

for painstaking care and experience. The

labor of measuring the stellar and comparison

lines on several hundred spectrograms is considerable and often hampers the progress of a

program.

The measurements were made originally

with the aid of a direct-vision microscope

mounted above a traveling micrometer, or, in

the Hartmann comparator, spectra were

matched against suitable "standards," again

with the aid of a microscope. The fatigue

accompanying this process was lessened, some

years ago, by the introduction of projection

methods (Petrie, 1937; Redman, 1939). A

further improvement was made at Victoria

when a new type of projection instrument

(Petrie and Girling, 1948) allowed one to

measure the Doppler shift directly upon the

projected image and, by virtually eliminating

the "reduction," gave a substantial increase in

efficiency. Indeed, so useful is this projection

comparator that the measurement and reduction of spectrograms is no longer a serious

obstacle to the completion of a heavy radialvelocity program.

The obvious advance to be made now is the

development of an automatic machine that
NO. 1 NEW HORIZONS IN ASTRONOMY 17

will remove the subjective errors and allow an

impersonal definition of the precise position of

a stellar line relative to the comparison spectrum.

A machine of this sort has been built by Johnson

(1949), and another of different type by Hossack (1953). In spite of these ingenious

machines there is still room for experimentation and development. The technical aids

already exist and it is fairly certain that the

present generation is the last that will be required to measure Doppler shifts with hand

and eye. A real triumph awaits some student

equipped with a sound knowledge of modern

electronic aids.

A word of caution may be permitted here.

An automatic machine does only half the job

if it merely presents one with the measured

positions of stellar and comparison lines. The

calculation of a radial velocity from the engine

readings requires nearly as much time as the

original measurement. The ideal machine,

therefore, will measure the shift of the spectral

line from its zero-velocity position and so allow

the rapid deduction of the radial velocity.

Wavelengths

The radial velocities deduced from the measured

positions of spectral lines are, of course, very

sensitive to the accuracy of our wavelengths.

We recall that in the photographic region an

error of 0.1 A is equivalent to 7 km/sec. A

firm knowledge of wavelengths is a prerequisite

to a significant knowledge of motions.

The practical difficulties are considerable.

Radial velocity programs usually force one to

employ a spectroscope of low dispersion and

purity; with such an instrument all spectral

features in types later than A 0 are blends of

two or more lines. The effective zero-velocity

wavelength of a feature depends upon the

relative intensities of the blending lines, the

spectrographic resolution, and the amount of

stellar rotation and other line-widening influences. Even in early-type spectra difficulties

sometimes arise because of blends, stellar

rotation, and Stark effect.

At Victoria a good deal of attention has been

devoted to the matter of effective wavelengths

and it is believed that satisfactory methods and

values have been established (Petrie, 1946,

1947, 1948, 1953: McDonald, 1948; Wright,

1952). The problems must be faced anew

with each new spectrograph and vigilance is

necessary to ensure that the deduced radial

velocities are, in fact, motions and not spurious

displacements arising from uncertain, or incorrect, wavelengths.

In spite of the work already done on this

subject, more is required. We do not, for

example, have a satisfactory set of standard

wavelengths for O-type spectra and cannot,

therefore, be entirely certain of their motions

in star "associations" nor of their apparent

"red shift" as members of galactic clusters.

And, at the other end of the temperature sequence, we lack effective wavelengths for lowdispersion spectra of N stars. This is of some

importance since we need to know the space

motions of a larger number of these (apparently)

faint objects.

Applications

The application of radial velocities leads us into

the subjects of stellar dynamics, galactic structure, binary stars, etc. These topics are

being discussed elsewhere in this compilation

and the writer will therefore merely point out

some of the possibilities awaiting the astronomer

who devotes himself to radial velocities. The

worker who undertakes to supply the fundamental observations of stellar motions can be

sure that his contribution is of permanent value

to astronomy. He must therefore be patient

and content to see his work gradually increase

as his programs and plans are completed.

The general problems of stellar motion remain with us: solar motion, star streaming,

galactic rotation, and mean velocities in different parts of the galaxy. The need is always

for more observations and observations of more

distant stars. For a general description of

present-day need one should refer to the report

of the International Astronomical Union Symposium on "Co-ordination of Galactic Research" (Blaauw, 1955).

The measurement of galactic rotation is of

prime importance. We require a more accurate

knowledge of Oort's constant A in the neighborhood of the sun. Current programs of B

stars at the Radcliffe and Dominion Astrophysical Observatories will supply about 1,000

new velocities, mostly within a distance of
18 SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS VOL.1

1,500 parsecs of the sun. These, together with

existing observations, should give a good

determination of the coefficient of differential

rotation in our part of the galaxy.

The more distant regions are quite imperfectly studied. With the recognition of

spiral arms about 2,000 parsecs away, it

becomes an exciting task to discover what

velocity effects exist at the edges of the arms

and between them. This means we must

observe large numbers of spiral-arm objects—

early B stars, Cepheids, and supergiants of

types B 8 to A 3. At the same time we must

measure the speeds of the interarm population

which will probably consist of stars of Population II together with a number of stars that

have wandered out of the spiral arms. Observations of the early-type stars will also

reveal the motions of the interstellar matter.

A problem of considerable interest is the

stellar velocity distribution perpendicular to the

galactic plane. This may be found from radial

velocities of stars near the galactic poles. The

observation of the brighter stars of types A 0-F 5

within 10° of the north Galactic Pole is now nearing completion at Victoria. A companion program

of fainter stars is proceeding at Mount Hamilton.

In this branch of stellar motion work much remains to be done. Extensive observations are

required among the later-type stars of the disk

population and, on the other hand, of the faint

blue stars of (probably) Population II. From

the velocities of the disk population we can

determine the density of matter in the galactic

plane region, in the neighborhood of the sun.

The above proposals involve observations of

many stars and are rather long-term projects,

but much may be done in other directions

where shorter programs will bring valuable

results. Some of these are suggested very

briefly in the following paragraphs.

Star clusters and associations.—These fascinating groups are of first importance in studies of

stellar luminosity and evolution. Radial-velocity observations are needed to verify cluster

membership and to calculate space motions.

Careful and extensive observations of the

"aggregates" and expanding associations must

be made before we understand the kinematical

properties of these groups. The study of

clusters and associations is sure to yield important and exciting information.

Stellar atmospheres.—Recent work on supergiant stars and special eclipsing stars, such as

31 Cygni and f Aurigae, shows how radialvelocity measures reveal the existence of atmospheric motions. This is a little-explored field

and requires study with highest possible dispersion. Comparisons between giant and mainsequence stars of the later spectral types will

be valuable in promoting knowledge of the

stability of stellar atmospheres.

Variable stars.—Experience has shown that

variable stars should be studied with simultaneous spectroscopic and photometric measures. The observations of Sanford (1952) and

Abt and Sanford (1954) on the Cepheid variable

W Virginis and of Struve (1954) on BW

Vulpeculae are good examples. The work to

be done here is limitless, considering the great

range of stellar variation. To be most fruitful,

the highest possible dispersion must be used.

Binary stars.—The problems of binary stars

are legion, but perhaps no other field offers more

challenges calculated to advance our knowledge

of stars.

Visual binaries should be observed continuously so that, ultimately, the combination of

spectroscopic and micrometric data will tell us

the dimensions and luminosities of the components. Similarly, eclipsing binaries should be

observed assiduously, especially those with

determinate light elements and those showing

the spectra of both components. Our knowledge of stellar masses, radii, and luminosities,

as well as further extensions and verification of

the mass-luminosity relation, must come from

this work. Radial velocity observers will find

it a rewarding task.

The bewildering array of spectroscopic anomalies exhibited by the close binaries has been

brought to our attention by Struve (1950).

The observation and interpretation of radialvelocity curves are an endless challenge. One

suspects that the most potent observational

aids will be required to advance our understanding of these spectacular phenomena.

Stellar satellites.—The discovery of planetary

systems is probably just within the powers of

the most precise radial-velocity measures that

can be made with existing equipment. This
NEW HORIZONS IN ASTRONOMY 19

has been shown by Struve (1952), who proposed

that a search be made. It may be that a major

astronomical discovery will result from the

meticulous observation of a number of our

nearby and bright stars.

References

ABT, H. A.

1954. Astrophys. Journ., Supplement, No. 3.

ADAMS, W. S.

1941. Astrophys. Journ., vol. 93, p. 11.

ARGYLE, P. E.

1955. Journ. Roy. Astron. Soc. Canada, vol. 49,

p. 19.

BLAATJW, A.

1955. Int. Astron. Union Symposium No. 1.

Cambridge University Press.

EBEBHARD, G.

1933. Handbuch der Astrophysik, vol. 1, pt. 1,

p. 299.

EDLEN, M. B.

1955. In Draft reports of the International Astronomical Union Dublin meeting, AugustSeptember 1955, p. 104. Cambridge

University Press.

FEHRENBACH, C.

1947. Ann. d'Astrophys., Paris, vol. 10, p. 306.

1948. Ann. d'Astrophys., Paris, vol. 11, p. 35.

HOSSACK, W. R.

1953. Journ. Roy. Astron. Soc. Canada, vol. 47,

p. 195.

JOHNSON, H. L.

1949. Astron. Journ., vol. 54, p. 190.

MCDONALD, J. K.

1948. Journ. Roy. Astron. Soc. Canada, vol. 42,

p. 220.

MERRILL, P. W.

1931. Astrophys. Journ., vol. 74, p. 188.

MlLLMAN, P.

1930. Circ. Harvard Astron. Obs., No. 357.

PETRIE, R. M.

1937. Journ. Roy. Astron. Soc. Canada, vol. 31,

p. 289.

1946. Journ. Roy. Astron. Soc. Canada, vol. 40,

p. 325.

1947. Journ. Roy. Astron. Soc. Canada, vol. 41,

p. 311.

1948. Journ. Roy. Astron. Soc. Canada, vol. 42,

p. 213.

1953. Publ. Dominion Astrophys. Obs., vol. 9,

p. 297.

PETRIE, R. M., AND GIRLING, S. S.

1948. Journ. Roy. Astron. Soc. Canada, vol. 42,

p. 226.

REDMAN, R. O.

1939. Monthly Notices Roy. Astron. Soc. Lon

don, vol. 99, p. 686.

SANFORD, R. F.

1952. Astrophys. Journ., vol. 116, p. 331.

SCHALEN, C.

1954. Arkiv fur Astronomi, vol. 1, p. 545.

STRUVE, O.

1950. Stellar evolution, chap. 3. Princeton Uni\

Press.

1952. Observatory, vol. 72, p. 199.

1954. Publ. Astron. Soc. Pacific, vol. 66, p. 329.

TREANOR, P. J.

1948. Monthly Notices Roy. Astron. Soc. London, vol. 108, p. 189.

WRIGHT, K. O.

1952. Publ. Dominion Astrophys. Obs., vol. 9,

p. 167.

Astrometry

By K. Aa. Strand*

Astrometry deals with the space-time behavior of the stars, and thus belongs to the

classical field of astronomical studies. While

the early investigations were directed mainly

towards establishing a suitable frame of reference for determining the complex motions of the

planets, studies of the positions and motions of

individual stars, as well as of the various stellar

systems, gradually developed.

If one examines the observational data accumulated in astrometry, the record is indeed impressive. During the past 150 years more than

370 catalogs of meridian observations of stars

have been published. We have monumental

astrographic catalogs like the Carte du Ciel, the

Astronomische Gesellschaft catalog (AGK2),

and the Yale zone catalogs, to mention only the

largest. The Yale General Catalogue of Stellar

Parallaxes lists the distances for nearly 6,000

stars measured since the beginning of this century, and the Burnham and Aitken double star

catalogs collect more than half a million micrometer measures, the result of a century's work.

Considering this impressive observational

record, it is not surprising that a person who

has not worked in astrometry usually assumes,

mistakenly, that most of the work has been

completed. Our knowledge of such properties

of the stars as their masses, luminosities, and

motions depends directly or indirectly upon the

available astrometric data, and any improvements to be made in these parameters will depend upon further work in the field.

What future work is needed in astrometry to

broaden our general knowledge of astronomy?

A whole series of problems confronts us. They

range in size from those that can be investigated

only by close cooperation among many observatories to those that can be solved by a few

astronomers in a^single observatory; and in

'^Director, Dearborn Observatory, Northwestern University, Evanston, HI.

381014—56 3

complexity the problems range from setting up

an accurate reference system of stars covering

the entire sky to determining the parallax of an

individual star.

A series of conferences held during the past

two years at Groningen, Evanston, Pulkovo,

and Brussels has resulted in general agreement

that, to solve many present-day problems, astronomers must undertake an international program of meridian circle observations. The program, starting in 1956, will consist of observations of nearly 21,000 stars brighter than 9.1

magnitude north of declination —5°. The total

number of observations, including those required of the fundamental stars, will amount to

260,000, since each star is to be observed 10

times. Twelve observatories will participate,

and by using timesaving devices in both the

observational and reduction phases of the program, they expect to complete it in 5 years.

American astronomers should note with some

pride that the U. S. Naval Observatory will

carry out a substantial part of the observations

and will also supply, on request, the reductions

to the day and to the equinox 1950 for all the

stars observed. The U. S. Naval Observatory

will also compile the final catalog of positions for

these stars.

What information will this large undertaking

give us ? It will provide a reference system of

stars with greater density and precision than

any previous system; and it will make possible

the solution of problems that could not be

solved without it, some of which may be summarized:

1. The program will provide a reference system for the AGK3, the second photographic observation of the stars in the Astronomische Gesellschaft Katalogs, which when compared with

the AGK2 will lead to accurate proper motions

for approximately 180,000 stars brighter than

11.5 photographic magnitude. Dr. Vyssotsky

21
22 SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS

of the Leander McCormick Observatory and his

associates will supply spectra for approximately

60,000 stars not previously classified.

2. About 11,500 of the stars in the reference

system will be used in the Russian program on

stellar proper motions relative to extragalactic

nebulae.

3. The proper motions determined from the

AGK2 and AGK3 will provide the systematic

corrections that are needed to make use of

many of the older star catalogs.

4. The results from the AGK3 will provide

the necessary data to reduce to absolute the

proper motions obtained from a reobservation of

the Carte du Ciel plates.

5. The stars from the AGK3 could be used as

reference stars for the Lick Observatory program, which expects to determine proper motions of stars with respect to extragalactic

objects.

It is gratifying to note that it is still possible

to undertake large-scale programs requiring

careful international cooperation, for which astronomy is conspicuous among the sciences.

However, it should be observed that the burden

of American participation in the AGK3 rests

solely upon the U. S. Naval Observatory.

Because the Lick program of deriving proper

motions of stars with respect to extragalactic

objects will yield important scientific results, it

should receive the support necessary for its completion. The Lick program will provide absolute proper motions of stars down to the 17.5

photographic magnitude, resulting in data that

can be expected to yield corrections to the

adopted values of precession and lead to new

determinations of solar motion and galactic rotation. The most distant stars in the Lick

survey will have distances of more than 2 kiloparsecs and thus reach beyond the region of the

maximum circular velocity, as indicated by the

Oort model. This study should settle the present discrepancy between the data on galactic

rotation, as obtained from proper motions and

from radial velocities.

If the Lick survey is to attain the needed

minimum accuracy, a time interval of approximately 20 years must elapse between first- and

second-epoch plates. This means that the complete repetition of the Lick plates will not begin

until 1967. How soon after that date the program will be completed depends upon the methods of measurement and of reduction. It is

quite clear, however, that the program is of such

magnitude that it requires the utilization of the

latest developments in automatic or semiautomatic equipment to complete it without prolonged delay.

Since the value of the Lick program would

be much enhanced if it could be extended to

the southern sky which is not observable at

the Lick Observatory, serious consideration

should be given to extending the program to

the Southern Hemisphere. In view of the

small number of observatories south of the

equator and the magnitude of the problems

that already confront them, this program probably should be undertaken by an observatory

in the Northern Hemisphere. It remains to

be investigated whether a telescope identical

with the Lick 20-inch Carnegie astrograph

should be constructed or whether, as suggested

by Dr. Luyten at the Meeting of the International Astronomical Union in Dublin, a

limited number of plates taken with the Bruce

telescope would serve the purpose.

Photographic astrometry with long-focus

refractors is traditionally a field in which

American astronomers have held the leadership; here, also, we face a series of problems

which must be solved for further progress in

astronomy.

It has long been known that the trigonometric

parallaxes are affected by systematic errors;

whether further statistical investigations of the

existing observations will provide additional

information on this problem remains doubtful.

The systematic errors have, of course, the most

serious effect upon the small parallaxes; but

the absolute magnitudes of giants and subdwarfs, which are scarce in the solar neighborhood, are therefore derived from a relatively

small number of small parallaxes and thus are

also affected by large systematic errors. To

improve this situation it is essential that many

of the small parallaxes be repeated, and that

the observational series for the individual

stars be increased from the 20 plates or less

used in the very early work to 50 plates or

even more. These determinations would add
NEW HORIZONS IN ASTRONOMY 23

to our knowledge of the absolute magnitude of

the stars just mentioned, and would also serve

as important statistical material for further

investigation of the systematic errors in trigonometric parallaxes.

The stars of low luminosity in the solar

neighborhood, discovered in the Luyten proper

motion survey, are another group for which

we need determinations of trigonometric parallaxes. The great majority of these stars,

late-type M-dwarfs, subdwarfs, and white

dwarfs, are considerably fainter than any we

can observe with the telescopes now being used

for parallax work. These telescopes cannot

reach stars much fainter than the 12 th magnitude, but most of the known white dwarfs

are fainter than the 14th. For this reason,

the commission on parallaxes and proper

motions of the International Astronomical

Union recommended at the recent meeting in

Dublin that observatories possessing large

reflectors seriously consider a program of

determination of trigonometric parallaxes of

faint stars.

Other aspects of the use of long-focus refractors are the positional study of the stars

within 10 parsecs by an extended series of

plates to discover perturbations for "single

stars," and the subsequent orbital analysis

of these perturbations. Quite recently such a

study led to the visual discovery of a binary

(Ross 614) in which the companion has only

8 percent of the sun's mass.

When we examine the relation of astrometry

to the double stars, we find a series of problems

which deserve the attention of modern astronomy. For the double stars within 20 parsecs,

it is possible to derive stellar masses with an

accuracy of 25 percent or better; to attain this

accuracy, however, visual and, whenever possible, photographic observations of the relative

motions are needed as well as extended series of

plates from which the mass ratios and parallaxes

can be derived with great precision. Much information in regard to stellar masses has been

obtained in this way over the past 15 years, but

much more is needed. This group of binaries

constitutes only about l/1000th of the approximately 40,000 known binary systems. To observe them all would clearly be an impossible

task; nevertheless, many of them must be observed if we wish reliable data on such questions

as statistical distribution of the orbital elements

in binaries and combination of spectra in physical pairs. These are only two of the questions

that have a bearing on stellar evolution and

which we cannot answer by studying the small

group of known binaries within 20 parsecs of

the sun.

Problems of membership and internal motions

in clusters will become even more important as

the time interval between early- and late-epoch

plates increases the accuracy. The early Carte

du Ciel plates and those taken with the large

refractors shortly after the beginning of this

century are especially important in that

respect.

One of the most recent additions to astrometry, the study of the age of the O-B associations from their expansion, deserves high priority, and it is to be hoped that the large meridian programs in connection with the AGK3 will

still permit the stars selected for this study to

be observed with utmost precision. A similar

program on Cepheid variables, to trace the positions of the spiral arms in our galactic system,

also deserves special attention.

Lack of manpower to cany out the many important problems just outlined remains a problem. We may summarize briefly what we require in instrumental developments to solve this

question. The staff of the U. S. Naval Observatory has, over a period of years, introduced

timesaving devices that have increased the

speed as well as the accuracy with which the

stars are observed, and because of the use of

IBM machines the reductions of the observations are no longer lagging behind as in the past.

In photographic astrometry the use of automatic devices in guiding, plate transport, and

timing has resulted in increased accuracy in

certain problems, but the real problem here is to

develop equipment that will effectively shorten

the time required for measurement and reduction of the photographic plates. One measuring

machine, allowing semiautomatic measurements

and operated in connection with IBM machines,

is already in use. Since the operating costs of

this system are prohibitively high for most
2 4 SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS VOL.!

American observatories, it seems desirable to tention. However, those outlined show that in

look into other recent technical developments spite of, or rather because of, the huge amount

which might aid in reducing the time spent at of work done in the past, there are still many

the measuring machine without requiring a interesting and important problems on which

highly expensive computing system. progress in astronomy depends; therefore, they

In this brief report it is not possible to men- deserve the attention of the active research

tion all phases of astrometry that deserve at- workers in astronomy.
Photoelectric Astronomy

By A. E. Whitford

 l

The past two decades have seen increasing

use of photoelectric methods in all branches of

astronomy where quantitative measurement of

radiation is important. This trend seems

destined to continue. In the years ahead the

photoelectric cathode may become the most

widely used radiation detector in the astronomical tool kit; indeed, it may supplant the photographic plate in the same way that the plate

supplanted the human eye 50 to 60 years ago.

A brief review of the links in the information

chain along which knowledge about the external

universe must pass will show the reasons for

believing that photoelectric methods may

assume a dominant place. The light beam from

a distant luminous source, which carries all the

information the astronomer will ever have

about that object, arrives at the surface of the

earth after some losses in the interstellar

medium and in the earth's atmosphere. The

atmosphere adds an unwanted background

glow, and, at times of bad seeing, diffuses the

beam somewhat. The potential information

in the light beam suffers a loss and degradation

before it enters the astronomer's instrument

that are for the most part beyond his control.

Some day the astronomer may get above the

atmosphere, but he will not get beyond the

interstellar medium. Meanwhile, he must

make sure that his earthbound instruments do

the best possible job of extracting all of the

information remaining in the light beam.

The astronomical telescope has long since

reached a state of perfection where the tranquility of the atmosphere, rather than the

quality of the optics, determines the image

quality. The astronomer may, in theory,

gather more information by increasing the

aperture, but the engineering difficulties and the

cost both rise steeply for any considerable increase in size over that of existing telescopes.

Wasbborn Observatory, University of Wisconsin, Madison, WJs

The final link in the information chain, the

radiation detector at the focus of the telescope,

is therefore the element which must be scrutinized for possible lack of efficiency. If the

efficiency is low, improvements can be counted

as equivalent to increasing the aperture of

existing telescopes.

The various practical detectors in actual use

can be evaluated by comparing them with a

hypothetical ideal detector that recorded the

arrival of each quantum, and did it with 100-

percent efficiency for all accessible wavelengths

of the astrophysical spectrum. The photographic plate cannot be given a very high mark

in such a comparison. Since it is a saturable,

nonlinear device, no unique quantum efficiency

can be assigned. But in the range of densities

favorable to contrast-judgments by the human

eye, only about 1 quantum in 1,000 is recorded

by producing a developable plate grain. Its

ability to integrate a long exposure is the reason

for its superiority over the eye in detecting

faint stars and recording low surface brightness;

the eye makes its decision in about % second.

Within that % second the eye, as shown by Rose

(1948), is a better detector than the photographic plate. The eye cannot, however,

deliver the plate's objective record of thousands

of details or the extremely precise positional

information needed for astrometry.

The cesium-antimony cathode is the one most

widely used in astronomical photometry. It

has a useful range of 3100A to 6000A. The

average quantum efficiency at the peak near

4400A is 12 percent, and exceptional tubes go

as high as 25 percent. There is thus a superiority of a factor of the order of 100 over the

photographic plate in the initial information

made available in a given time from a given

angular area on the sky. As presently used,

only one chosen element or area may be measured at a time; if photometric information on

25
26 SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS VOL.1

many objects in one field is needed the photographic plate is a great timesaver. But the

superior quantum efficiency of the photoelectric cathode is the basic reason why it

seems destined to play an increasing role in

measurements of astronomical radiation.

In order for the favorable quantum yield to

be useful, there must be no instrumental background or amplifier noise large enough to

degrade appreciably the information contained

in the released photoelectrons. Existing techniques succeed quite well in approaching this

ideal and further efforts can be expected to

produce only small improvement. There is not

much to choose between any of the three

currently used methods of registering the output

of a phototube: current measurement with a

load resistor, charge integration with a storage

condenser, and charge integration by pulse

counting. At the lowest intensities integration

methods will probably find increasing use.

The future uses of the photoelectric method

in astronomical photometry are of course not

entirely predictable. The only really safe

prediction is that new and ingenious applications of the method, not now foreseen, will

surely be devised. A projection of the presently successful techniques suggests activity

in three main areas. In general, these exploit

the basic advantages of the photoelectric

process: high quantum yield, linear response,

and wide spectral range.

Intensity comparisons: Magnitude systems

The range of intensity between the sun and

the faintest stars yet detected is of the order

of 1020. Important astronomical conclusions,

such as the cosmic distance scale, depend on

the precision with which this range is covered.

The superior accuracy of the photoelectric

method rests in part on its higher quantum

yield, in part on its linearity. Elimination of

the step calibration needed by the nonlinear

detectors such as the photographic plate not

only saves a great deal of time but also removes

a fertile source of accumulated error. An

intensity range of 10* to 106 can be covered

with one phototube and a single optical system.

The task of calibrating magnitude systems

is already largely a photoelectric province.

Photographic interpolation between photoelectrically established standards in the field

of a star group such as a globular cluster

(e. g., as by Arp, Baum, and Sandage, 1953)

will continue to be used where photoelectric

measurements of one faint star at a time would

be too time-consuming. But for transfer of

standards from Selected Areas, photoelectric

methods appear to be needed for good accuracy.

The time is rapidly approaching when every

observatory which does any photometry at

all will think of the plateholder and the photoelectric photometer as routine interchangeable

attachments to the telescope, each to be used

as occasion demands.

Since photoelectric methods require the use

of the telescope for only one object at a time,

and the reductions are also quite time-consuming, magnitude measurements have thus far

been restricted to specific classes of stars, such

as B-stars and Cepheid variables, and to the

establishment of standards in Selected Areas

or star groups under special study. A determined effort might provide instrumental and

computing aids which would make mass production of magnitudes more economical. These

might include better coordinate mechanisms

for rapid identification and centering of stars,

automatic subtraction of sky readings, automatic computation of extinction and correction

of the star signal for it, a logarithmic amplifier

to give direct reading magnitudes, and direct

processing of telescope output onto data storage

records such as punched cards. Then it might

not be too enormous a task to do what Irwin

(1955) suggests, and undertake to provide

accurate photometric data for all the stars,

say, in the Henry Draper Catalogue. A massive

attack by one observatory could accomplish

it; it would be a public-spirited endeavor, to

be viewed as a repayment of the scientific debt

owed to previous generations for the patient

determination and compilation of star positions,

spectral types, radial velocities, and proper

motions. It would at least correct the present

incongruous situation where the simplest and

most fundamental datum about a star, its

brightness, is often the least reliable entry in

the catalog.
NEW HORIZONS IN ASTRONOMY 27

Color systems and spectrophotometry

The advantage of the photoelectric cathode

in color studies stems partly from its wide

spectral range, partly from its linearity. The

simple ratio of the intensity through two color

filters gives a good color index without the need

of establishing two magnitude scales as in the

photographic method. Because the photoelectric color index comes from a differential

measurement, the internal accuracy is higher

than for the magnitudes themselves. A twocolor index provides the minimum amount of

information about a star. The wide spectral

range of a photoelectric cathode permits three

or even five narrow-range colors with a cesiumantimony cathode. The range may be extended

to 10,000A or more with a cesium-oxide cathode

and the number of color filters extended to six

or seven. The next step is to scan the whole

range with an exit slit on a spectrograph. The

observing time per star goes up because of the

limitation that only one element of the spectrum

can be observed at one time. High-resolution

photoelectric scans of short portions of the solar

spectrum have been made, and similar scans will

undoubtedly be attempted for the brightest

stars. Over the whole range, however, a resolving power of about 350 gives a reasonable

observing time per star and adequate delineation

of the main features of the continuous spectrum.

The color of a star is already recognized as a

necessary part of a photometric determination

of its distance in order to evaluate the interstellar reddening. For early-type stars, at least,

three colors on the U-B-V system of Johnson

and Morgan can give a measure of reddening

when the spectrum is not known. Extension of

this powerful method to later-type stars, and

to the longer wavelength portions of the spectrum where the interstellar absorption is less,

would be a great help in galactic and extragalactic distance determinations.

Determination of the absolute energy distribution in normal Type I stars will be worthy of

considerably more work over the whole range

that can be covered with existing photoelectric

cathodes. From such studies can come better

model atmospheres, and from these, in turn,

higher accuracy in quantitative studies of nuclear abundances from line strengths. Comparison of these "normal" stars with Type II

stars and peculiar stars by multicolor or scanning methods should be most informative. Examples would be subdwarfs, metallic line stars,

and pulsating stars of all kinds. Differences of

composition or structure may be expected to

leave clues in the continuous spectrum; accurate photometry will be needed.

Composite sources are another fruitful field

for multicolor studies. Globular clusters and

galaxies are examples. The resolving power in

distinguishing between various syntheses of the

over-all energy distribution from hypothetical

mixtures of known stellar types may never be

very high. But a combination of spectra, of

counts of brightest stars of identifiable types,

and of over-all energy distribution may lead to

fairly definite conclusions. The diffuse nature

of these objects precludes anything but low

resolution spectrophotometry by the scanning

method, and for faint objects multicolor filter

photometry may be the best that can be done.

Their low contrast with the sky can be handled

cleanly by the differential subtraction methods

so easily applied in photoelectric photometry;

and photoelectric accuracy will be needed.

Photometry of distant galaxies for diameter and

magnitude (as suggested by Baum, 1953) and

for energy distribution (Whitford, 1953) will

provide evidence of the state of affairs at a

much earlier era in the evolution of the universe,

and will thus be of interest to those concerned

with cosmogony.

Differential measurements

Light curves of eclipsing binaries provided

one of the early demonstrations of the precision

attainable by photoelectric methods. The

accuracy was due partly to the linear response,

and partly to the differential comparison

of variable and comparison star, thus minimizing variable factors in sky and apparatus.

The fundamental information about stars

derivable from precision light curves of eclipsing

variables will continue to offer attractive possibilities to the photoelectric observer with a

small- or medium-sized telescope. For reconnaissance purposes other methods will continue

to be of value, but the precision needed for

deriving properties like limb darkening requires

photoelectric accuracy.
2S SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS VOL.1

Linear response is a great aid in all manner

of differential measurements because simple

subtraction is possible. Alternation of fields

once a minute or oftener is one way of doing

this. The flicker or a. c. method is another,

and it may give a better elimination of unwanted

extraneous factors. One example is the measurement of very faint stars or galaxies which

add only a few percent to the sky radiation

that comes through the smallest usable focalplane diaphragm. Nebulosity of low-surface

brightness would be another example. Extension of flicker methods to this type of

measurement would be a logical development.

Wherever it is possible to "put a tag" on

the wanted signal, it can usually be lifted out

of an overwhelming background by differential

methods, often involving a. c. amplification.

One example would be the polarization of

starlight (Hiltner, 1949, and Hall, 1950);

another would be measurement of solar magnetic

fields (Babcock, 1953) where the modulating

element again involves polarization. These

flicker methods have been highly developed in

radio astronomy where the signal may be 100

to 1,000 times smaller than the unavoidable

noise. Further use of the technique in photoelectric astronomy seems likely.

Possible advances in technique

If photoelectric astronomy continues under

the restriction of one picture element at a time,

what developments might extend its range of

usefulness ? At the head of the list would come

a more sensitive cathode for the red and infrared where the quantum efficiency is now 20

times poorer than for the blue. This region

of the spectrum may assume increasing importance for two reasons: Interstellar absorption is greatly reduced, and discovery level

or penetrating power correspondingly increased;

and at very large red shifts, extragalactic systems

carry not only their energy maximum but also

all the parts of the spectrum ordinarily studied

into the red region. A high-efficiency photoelectric detector would greatly facilitate investigations in these areas. A new tri-alkali

cathode (Sommer, 1955) promises considerable

improvement in the red as far as 7500A.

Germanium phototransistors may be developed

sufficiently to fill the need as far as 16,O0OA.

Some reduction in observing time can be

realized by using two or more phototubes, fed

by beam-splitters, or by several different outputs of a spectrograph. The most exciting

development in photoelectric astronomy, however, would be an image tube that would

preserve the high quantum yield of a cathode,

while responding simultaneously yet separately

to all the bits of information in the various

area elements of the image plane. Such a

development seems to be a definite possibility.

The techniques being explored, as explained

by Baum, have been reported by Struve

(1955). Work is actively in progress in

France, England, and the United States.

In its simplest form, that first introduced by

Lallemand, the image tube consists of an electron lens system for focusing the accelerated

photoelectrons in an imago on a photographic

plate at the end of the tube opposite to the

cathode. Every electron sensitizes a plate

grain, thus preserving the high quantum efficiency of the cathode. The same number of

developable grains is achieved in a much shorter

time than would be needed in direct exposure of

the plate to light. A fine-grain plate may be

used without loss of speed, thus increasing the

number of bits of information in a given image

area before saturation sets in. As yet, exposure

times adequate for showing an intensity level

as low as the night sky have not been possible,

and the anticipated gain in threshold detection

level has therefore not been realized.

A proposed alternative scheme involves

storage of the released charge on a mosaic, as

in the image orthicon. This is later wiped off

by a scanning electron beam, and the effect

on the beam is the signal from that element.

While a great deal of hard work remains,

there appear to be no insuperable difficulties

in the way of final success. If such a tube

comes, what will be the first applications?

Reduction of prohibitively long exposure times

in photographic spectroscopy would be one of

the first. All-night exposures would be reduced to a few minutes. Spectra of faint

galaxies, so important in considering cosmological questions, would become fairly routine. On direct pictures of the sky a detection

limit 2.5 magnitudes fainter should result,

because 100 times as many countable events
NEW HORIZONS IN ASTRONOMY 29

go into the decision as to whether one picture

element (say the size of the seeing tremor disk)

has received a statistically significant excess

of photons from a star during the exposure

•time.

While high quantum efficiency and accurate

imaging can be realized in an image tube, it

may be doubted whether the linearity and other

photometric advantages of a single phototube

can be preserved through the whole process.

Nevertheless a successful image tube would

bring many observations, now utterly unattainable, into the realm of the possible. In

the rich harvest that would follow its introduction, there would almost certainly be some

unexpected discoveries.

Conclusion

As one looks ahead, it seems clear that the

next decade will see increasing use of the photoelectric process in astronomical photometry.

Wherever two quantities of radiation have to

be compared, the photoelectric method will be

the preferred method if time permits. Should

the image tube come into common use, a revolution in astronomical observing methods may

ensue. It is not impossible that the whole

field of the detection and measurement of

radiation may pass from the present state of

low efficiency to something approaching the

ideal state where every quantum is counted.

It would then resemble the present state of

optics as an astronomical tool: no revolutionary

improvements could be expected. But even

then, as now, astronomical advances will

depend quite as much on the skill and ingenuity

of the astronomer as on the tools he has at

his disposal.

References

ABP, H. C , BATTM, W. A., AND SAN DAG E, A. R.

1953. Astron. Journ., vol. 58, p. 4.

BABCOCK, H. W.

1953. Astrophys. Journ., vol. 118, p. 387.

BAUM, W. A.

1953. Astron. Journ., vol. 58, p. 211 (abstract).

HALL, J. S.

1950. Publ. U. S. Naval Obs., No. 17, pt. 1.

HlLTNER, W. A.

1949. Astrophys. Journ., vol. 109, p. 471.

IBWIN, J. B.

1955. Sky and Telescope, vol. 14, p. 132.

ROSE, A.

1948. In Marton, ed., Advances in electronics,

ch. 3. Academic Press, New York.

SOMMER, A. H.

1955. Rev. Sci. Instr., vol. 26, p. 725.

STBUVE, O.

1955. Sky and Telescope, vol. 14, p. 224.

WHITFOBD, A. E.

1953. Astron. Journ., vol. 58, p. 49 (abstract).

881014—56 4

Interferometers in Radio Astronomy

By B. F. Burke * and M. A. Tuve

 x

Only a faint glimmer of an idea yet exists as

to how astronomical objects can radiate, apparently with a fair degree of conversion efficiency,

electromagnetic waves some tens of centimeters

or even meters in length. The discovery phase

is probably over; only a few such objects are

known, but these are enough to face us unequivocally with the problem. "Plasma oscillation" is an easy phrase, and where the electron

densities are high, as in the solar atmosphere,

these words seem to offer a plausible explanation, but when one is confronted with the

problems of boundary conditions and the mechanism for escape of the radiation to free space

the plausibility disappears. Additional difficulties arise for the low apparent electron densities

in gas clouds. Strong and perhaps turbulent

magnetic fields combined with violent streaming motions, perhaps even with shock waves,

may resuscitate the hypothesis, but theory has

yet to rescue the observer in this field.

Progress toward solution of this problem may

still lie in the hands of the observer, however,

as there exist only a dozen or so sources which

have been specifically located and for which

"color" information (in the radio spectrum), for

example, has begun to be available. When at

least several hundred sources, thermal and nonthermal in their radiation characteristics, have

been cataloged and their radiation properties

described in some detail, the puzzles may diminish. Large arrays of antennas, mainly

dipoles or helices, connected and used as interferometric devices, are the principal technical

basis for obtaining observations at the longer

wavelengths. Parabolic reflectors of any reasonable size cannot be expected to give the

necessary precision of> location in the sky; at

20-cm wavelength even a 500-foot parabola

1

 Department of Terrestrial Magnetism, Carnegie Institution of

Washington, Washington, D. C.

would still cover an area 5 minutes of arc in

diameter. Arrays of great size and interferometer spacings of as much as a mile are simple

by comparison to such a dish.

It should not be surprising if a new branch of

science raises difficulties in its early days, and

radio astronomy has not been disappointing in

this respect. In the past few years, this young

branch of astronomy has provided a large

variety of both physical and astronomical

problems. A striking feature of the earliest

investigations has been that most of the intrinsically strong sources of radio noise are

relatively insignificant optical objects, a circumstance that adds to the difficulty of examining

in detail their physical behavior. The handful

of optically identified discrete sources all exhibit

one feature in common, however, which provides some indication of the nature (and difficulty) of the problem, for in all cases we seem

to be dealing with highly turbulent gas whose

constituents are in high states of ionization.

Other general or special requirements may

exist, and they will have an important influence

on any theory of how such a tenuous but

violently agitated gas becomes a relatively efficient emitter of radio noise. For example, the

polarization of the light from the Crab Nebula

(pi. 1) provides evidence for the existence of a

magnetic field and large quantities of highly

energetic electrons, which also could be efficient

radiators in the radio spectrum, but whether

this is a general characteristic of all radio

sources or a peculiarity of this one object only

the future will tell. It appears certain, however, that the necessary highly turbulent gas

can be produced in a variety of ways. In our

own galaxy, cataclysms such as supernova

explosions are known means of providing a

turbulent medium that is capable of strong

radio emission; but although all discrete galactic

31
32 SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS VOL.1

sources may possibly be supernova remnants,

many of them may also represent an entirely

new class of objects.

The extragalactic sources, which are intrinsically more powerful emitters of radio noise,

require a correspondingly more violent origin.

The collision of two galaxies seems to be a

sufficiently drastic occasion, but there are peculiar galaxies such as M87 which evidently require some other means of excitation.

The existence of two different populations of

radio sources raises the important question of

how to distinguish between them, and how to

establish, if possible, a distance scale. The

only reliable method at present seems to be identification with a visual object, but as time goes

on this method must prove increasingly difficult.

Even now it is clear that if a pair of colliding

galaxies as violent as the Cygnus source exists,

but 10 times farther away, we have little hope

of observing it visually even though it would

still be an observable radio source. Similarly,

many galactic objects, conspicuous in the radio

region, are undoubtedly so obscured by dust

that optical observation will be difficult. In

both cases, however, a distance scale would certainly bring to light interesting new problems.

For example, the distribution in space of extragalactic radio sources may provide a clue to important cosmologies! problems, since it is probable that many observable radio sources are

situated at great distances, approaching the

limits of the visible universe. So far, only statistical methods have been employed, and at

present there appears to be a serious discrepancy

between observations. The results of the Cambridge interferometer radio survey indicate an

unusually large number of faint sources, implying an actual increase in apparent spatial density

of radio sources in the universe at great distances. On the other hand, preliminary observations by Australian radio astronomers show

a number-magnitude relation that implies a

nearly uniform space density of radio sources.

The resolution of this direct observational conflict is certainly one of the most pressing problems in radio astronomy today. The solution

may lie in an understanding of the techniques

involved, for both surveys utilized interferometric techniques but in radically different

forms.

All the sources we have been considering are

characterized by their nonthermal nature, since

nonequilibrium processes appear to be involved.

An equivalent black body in order to exhibit an

equal surface "brightness" in the radio region

not only would have to be at an extraordinarily

high temperature but also would have to have

an equivalent temperature that rises with increasing wavelength. In the meter-wave region, where many sources are sufficiently intense to be seen with relatively simple equipment, they appear superimposed on a smooth

background of radio noise, most of it evidently

of galactic origin.

Interferometric techniques are particularly

valuable in removing the effects of the background, "labeling," as it were, the signals originating in discrete sources. Figure 1 illustrates

the principles of the phase-switching interferometer first described by Ryle (1952). In the

upper drawing, two antennas are arranged

symmetrically to form the radio analog of

the Michelson interferometer. If a radio source

is directly overhead, or is off-axis by an amount

sufficient to result in a path difference of

an integral number of wavelengths, the two

signals constructively interfere. Therefore, by

switching the phase of one antenna at a regular

rate the signal from a discrete radio source is

modulated, while the slowly varying background is not, and we1

 can distinguish the noise

from such sources by looking for a modulated

component in the total signal received. The

system has other advantages as well. If a

source is at a half-power point of the interference pattern, the signal received will be the

same whether a half wavelength has been inserted or not, and consequently at these positions the signal from a discrete source is unmodulated. As a result, position measurements

can be made by a null method, with much

greater precision than is possible with a pencil

beam. Furthermore, the relative amplitude of

the interferometer signal as a function of antenna spacing gives a measure of the angular

size of the source, effectively constructing the

Fourier transform, just as in the optical Michelson interferometer. It is possible, therefore, to

obtain information with relatively small antennas that otherwise could be obtained only

with a single antenna of prohibitively large size.
BURKE AND TUVE PLATE 1

The Crab Nebula. Photograph from the Mount Wilson and Palomar Observatories

(200-inch photograph).

NEW HORIZONS IN ASTRONOMY 33

I I I I •>v

ANTENNA

A

CABLE

IRECEIVERI

CABLE ANTENNA

B

DIRECTIONAL RESPONSE WITH EQUAL CABLE

LENGTHS; SIGNALS FROM OVERHEAD SOURCE ADD

111 1

L _

1 1 < •

at

o

I RECEIVER! -g- X

HALF WAVE DIFFERENCE IN CABLE LENGTHS:

SIGNALS FROM OVERHEAD SOURCE CANCEL

FIGURE 1.—Phase-switching interferometer.

There is danger in the method, however,

which must be guarded against; if more than

one source is being received, the interferometer

will measure the combined effect of the two,

and, unless one reconstructs the Fourier transform by going to sufficiently large spacing, a

single source will appear to be observed where

in reality there are two. The effective gain of

the interferometer is the geometric mean of

the gains of the two antennas, and at first sight

it would appear that if we simply make the

antennas larger, and thus see fewer sources per

beam width, this difficulty would be circumvented. It is true that the confusion of stronger

sources may be resolved by this technique, but

the increased antenna gain then makes it possible

to detect a still larger number of sources per

beam width, and since one naturally wishes

to examine these, also, multiple spacings are

still required. At the longer wavelengths,

particularly, where a pencil beam of the order

of a minute of arc or less requires an array of

prohibitively large size, interferometric techniques are indispensable, but they must be

applied with caution and full knowledge of the

difficulties involved.

The study of the sun is an interesting special

case, for the radio emission at longer wavelengths provides a useful tool for probing the

corona; here interferometers have already

proved their value, providing us with a qualitative picture of the sun at a number of different wavelengths. The detailed appearance of

the sun, particularly at the longer wavelengths

(greater than 4 m) needs much more study,

and there is little doubt that interferometric

techniques must prove necessary.

It was remarked earlier that the effective

sensitivity pattern of a pair of antennas used

as an interferometer is equal to the geometric

mean of the gains of the individual antennas.

This fact is utilized in the special case of the Mills

Cross (Mills and Little, 1953), which, in reality,

is a zero-spacing, phase-switched interferometer

utilizing two linear arrays at right angles. Figure 2 gives an illustration of the principle.

Each array has a fan-shaped sensitivity pattern,

but when connected as a phase-switched interferometer the effective sensitivity is the geometric mean of the two antenna patterns, which

is a pencil-beam whose resolving power is equivalent to a square array containing many more
34 SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS VOL.!

MAIN LOBE

OF X ^

PRINCIPAL LOBE WHEN

PATTERNS ARE MULTIPLIED

MAIN LOBE

OF "B"

ARRAY 7C ARRAY

OUTPUTS TO RECEIVER

FIGURE 2.—Antenna diagram of the Mills Cross interferometer.

elements. At the longer wavelengths this provides us with a technique for producing a pencilbeam with simplified arrays, and the future will

undoubtedly see arrays based on this principle

but of far larger dimensions than those in use

today, since a pencil-beam is unquestionably

the best way to investigate extended sources.

As in most experimental and observational

fields, the advantages of this method are accompanied by disadvantages as well, for this

system is troubled by spurious responses arising from its unusual diffraction pattern. It

is possible for a strong source, passing well

outside of the main pencil beam, to appear as

a weak source, and great care must be exercised

to avoid, if possible, or at least to identify such

spurious responses. (See fig. 3.)

The longer wavelengths, of the order of a

meter or more, have been emphasized since

the observation of a large number of sources

at these wavelengths can probably be accomplished more conveniently by interferometric

techniques than by an equivalent array, while

the measurement of angular size of small

sources is probably possible by no other means.

The phase-switching technique has the further

advantage that the noise originating in discrete

sources is "labeled" by modulation, and if, in

the future, it is possible to investigate the

intrinsic properties of source noise this feature

may be invaluable in distinguishing the desired

signal from the galactic noise background.

All present interferometric systems have weaknesses, however, which must be appreciated

if the important work on weak radio sources

is to have reliable significance.
NEW HORIZONS IN ASTRONOMY 35

PATH

OF

STAR

FIGURE 3.—Detailed antenna pattern of a Mills Cross interferometer.
36 SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS VOL. 1

In our concern with equipment it is important

that we not lose sight of the larger issues. The

cosmological problem has been outlined, but other

important questions also exist. What is the

cause of radiation in the case of those galaxies

that are not, apparently, colliding with anything? What is the significance of the distribution of radio sources in our own galaxy?

And, what are the physical conditions necessary

for the emission of radio noise from large,

turbulent gas clouds? Are magnetic fields

essential? The answers to these questions,

and others as well, depend on an increase in

the information concerning discrete radio

sources. Larger interferometers will certainly

play a role in the gathering of this information.

References

MILLS, B. Y., AND LITTLE, A. G.

1953. Australian Journ. Phya., vol. 6, p. 272.

RTLB, M.

1952. Proc. Roy. Soc. London, ser. A, vol. 211,

p. 351.
Related Sciences

Rocketry

By R. Tousey *

For many years astronomers the world over

have dreamed that sometime it might become

possible to discover the nature of the extreme

ultraviolet and X-ray spectrum of the sun and

the stars. Research in this field became

possible at the conclusion of World War II

as a result of the development in Germany of

the V-2 rocket, which was capable of carrying

scientific equipment to altitudes far above the

main absorbing regions of the earth's atmosphere. The V-2 rocket and the American

Viking and Aerobee rockets which followed

have opened the field of rocket astronomy.

The cost of this research is large, but the stakes

are very high, for it seems probable that the

key to many unsolved astrophysical problems

may be found in the nature of the spectrum of

the radiation of very short wavelengths emitted

by the sun and the stars.

The difficulties of conducting astronomical

research from rockets are many. The space

available for instrumentation in all but the

largest rockets is rarely more than 10 cubic

feet, and the payload is of the order of only a

few hundred pounds. The time available for

observation is 2 to 5 minutes. It is difficult

to keep instruments pointed at the celestial

object under study since rockets roll and turn

in an unpredictable fashion. Flight preparations are time consuming and complicated, so

that the firing time is often determined by

circumstances other than those that are optimum for the experiment.

In spite of the vicissitudes of research from

rockets, during the past several years great

advances have been made toward the solution

of the problems involved in rocketry; yet, there

is much still to be desired. Astronomers are

most anxious that the work be carried ahead as

vigorously as possible.

1

 U. S. Naval Research Laboratory, Washington, D. C.

Improvements needed in rocketry

Although the V-2 rocket was capable of

carrying a tremendous payload to an altitude

of 100 miles, which is adequate for many

experiments, it was extremely complicated,

difficult to launch, expensive, and more often

than not it misfired. Two American rockets

were designed to replace the V-2 for upper

atmosphere research: the Aerobee, a small,

comparatively inexpensive rocket with a peak

altitude of about 120 km; and the Viking,

similar to the V-2 hi complexity and size,

capable of 200-km altitude, and stabilized

during flight. A number of Vikings were

flown with successful experiments, but the

trend toward economy has forced scientists

to use and improve the Aerobee rocket for

research.

The Aerobee uses an acid-aniline fuel, is

launched with a booster from a tower, and

costs at present about $30,000. This rocket is

quite satisfactory for many experiments. It

can be launched at sea, though not as easily

as is desirable. The early type of Aerobee did

not reach altitudes sufficient for many experiments. The Aerobee-Hi, now in production,

will ascend to 200 to 300 km with payloads between 150 and250 lbs.; thus,it has performance

characteristics that are satisfactory for almost

all types of research. Although it is a much

simpler and cheaper rocket than the V-2 or

Viking, the Aerobee is still too expensive to be

used as frequently as would be desired in

astronomical research.

The preparation of a rocket and the launching are complicated. Furthermore, at land

bases such as the White Sands Proving Grounds,

the scientist is limited to firing at a particular

tune, when the range is ready and clear and the

winds are favorable, even though he may wish

to fire, say, during the onset of a solar flare.

39
40 SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS VOL.1

The improvement most needed in rocketry,

therefore, is a really inexpensive rocket that

can be made ready and flown without elaborate

preparations and precautions. Recent advances encourage hope that such a rocket,

employing a solid fuel propellant, can be developed in the next few years. A rocket that will

carry 100 pounds to 300 km and can be launched

by a few persons, perhaps by the scientists

themselves, at a cost of perhaps $5,000 no longer

seems unattainable. Such a rocket might best

be launched at sea, or at least flown out to sea

from shore. This would avoid complicated

and restricting safety requirements. Sea recovery methods would have to be worked out,

since most astronomical experiments are photographic. This involves the sealing and flotation of the section containing the instrumentation and the development of a system for locating the floating section.

For flights over land, recovery of photographic

film may be achieved simply through the use

of a strong casette, but ordinarily the equipment is completely ruined. Recent advances

in parachute techniques, however, make it

possible now to salvage the equipment for

further study or reuse. The cost of the parachute assembly must be greatly reduced,

perhaps to the neighborhood of $500 to $1,000,

before its general use will be warranted economically.

One of the most difficult and important

problems in rocket astronomy is that of keeping

the equipment accurately pointed at the celestial object being studied. Although it may be

possible to stabilize a rocket in flight, the precision of equipment stabilization required for

many experiments can be far more easily and

cheaply provided by stabilizing only the

instrument. So far, attempts at pointing have

been limited to the sun. The most successful

of several "sun-followers" is the "Biaxial

Pointing Control" developed by the University

of Colorado under Air Force sponsorship and

described by Stacey, Stith, Nidey, and Pietenpol

(1954). This device kept a spectrograph flown

in a U. S. Navy Aerobee on February 22, 1955,

pointed at the sun within an average error of

± 1 minute of arc during a total exposure time

of 2 minutes. The cost of this equipment is

high, therefore safe recovery is mandatory if

many experiments are to be conducted.

Further development of pointing controls in

the direction of lower cost and greater loadcarrying ability and simplicity is very important

for rocket astronomy. With such controls the

study of the bright stars from rockets may

become possible. This will probably require a

pointing control that operates from the moon,

or the sun, and the earth's magnetic field, or

possibly from simply the moon and the sun;

the star itself may provide the final signal for

precise pointing.

The fields of rocket astronomy

The experiments already conducted from rockets cover many fields of research in astronomy and the physics of the upper atmosphere.

The work already accomplished has been most

recently summarized by a collection of papers

presented at the Gassiot Committee meetings

(Boyd and Seaton, 1954). Many fascinating

new experiments have been proposed. In this

section an attempt is made to cover some of the

most interesting.

The solar spectrum.—The sun is certainly the

most important and the least difficult celestial

object to study. From the rocket work conducted by the Naval Research Laboratory

(Tousey, 1953a, 1953b; Wilson et al., 1954;

Johnson et al., 1955), by the University of

Colorado (Rense, 1953), and by the Air Force

Cambridge Research Center (Jursa et al., 1955),

its ultraviolet spectrum is now known in considerable detail to 977 A. Friedman (1955) and

coworkers (Havens et al., 1955; Friedman and

Chubb, 1955; Byram et al., 1955), using photon

counters, have measured the sun's X-ray

emission in several wavelength bands between

100 A and 5 A, and also have obtained much

information about the Lyman alpha line of

hydrogen and the dissociation of oxygen in the

upper atmosphere.

Down to 2085 A the ultraviolet solar spectrum resembles, generally, the visible spectrum

in having many deep Fraunhofer lines. The

intensity continues to fall below that of a 6,000°

K black body, and at 2100 A has reached a

brightness temperature of approximately 5000°

K. The photon counter measurements indicate
NEW HORIZONS IN ASTRONOMY

4,050° K at 1450 A and 3,900° K at 1250 A.

At these wavelengths the continuous spectrum

must come from the very coolest layers of the

sun's atmosphere. At 2085 A strong continuous absorption of unknown origin sets in;

although nearly obliterated, the Fraunhofer

lines can still be seen, the lowest so far identified

being at 1700 A. The continuum itself has

been photographed to 1550 A.

Below 1650 A the solar spectrum consists, in

the mam, of emission lines. So far, 32 emission

lines have been identified. These include Lyman alpha and beta of hydrogen, the two

resonance lines of O VI at 1032 A and 1038 A,

one resonance line of N V at 1239 A, the line

of He II at 1640 A corresponding to Ha, and

various (mostly) resonance lines of silicon and

carbon up to three times ionized. The intensity

of Lyman alpha has been measured on several

occasions, both photographically and with

photon counters, with results varying from 0.1

to 1 or more ergs/cm2/sec, suggesting that the

emission in this line may be quite variable,

depending on solar conditions. The soft X-ray

intensities measured with photon counters

appear to be roughly those expected from a

corona at 750,000° K. They have been shown

by Havens, Friedman, and Hulburt (1955) to

be quite sufficient to produce the E-layer of the

ionosphere.

There remains a great deal of fascinating research on the sun to be done from rockets.

In this field only the surface has been scratched.

While the data already obtained bear closely

on the structure of the sun, much more needs to

be learned to make possible the formulation of

a complete picture of the solar atmosphere.

The solar spectrum from 977 A to 100 A is still

completely unknown, and from 100 A to 5 A it

has been measured only in broad chunks. In

this spectral range lie the resonance lines of most

of the highly stripped atoms present in the

corona. It is now possible to study this region. The requirements are a good pointing

control and a rocket that will spend several

minutes above 200 km. A scanning spectrograph with a photon counter detector and

telemetering may be the simplest approach, at

least to obtain information without great

resolution. In the end, however, spectrum

photographs should be obtained.

Special studies of the sun.—The solar spectrum

is most easily studied by using the light coming

from a considerable area of the solar disk.

It is important, however, to learn exactly

where on the disk the various emissions originate, and how their intensities depend on the

height in the solar atmosphere. Several approaches to this problem have been proposed.

Some information can be obtained with a

stigmatic spectrograph on whose slit the sun's

image is held stationary to within 1 minute

of arc with a pointing control. Spectra

obtained in this fashion by Johnson et al.

(1955) show, for example, that the O VI line

at 1032 A has a greater ratio of limb to center

of the disk intensity than does Lyman beta at

1026 A, as would be expected. More measurements of this sort would make it possible to

improve greatly the model of the chromosphere and corona.

Ideally one would like to fly a spectroheliograph and to measure the intensity distribution over the disk and corona of all lines in

the vacuum ultraviolet. For the Lyman alpha

line, which has very high intensity, such an

instrument is feasible. A possible design was

published by Behring et al. (1954). Several

Lyman alpha spectroheliographs or cameras

will be flown in the next year or two. Following

this, weaker lines of higher excitation energy

should be investigated, including the ultimate

lines of He I and II at 584 A and 304 A,

respectively.

Another method of studying the distribution

of Lyman alpha over the disk, suggested by

Friedman, employs a photon counter of a

narrow field of view, sensitive to this line only.

By using a television scan of about 20 lines

resolution, a crude picture could be transmitted

to ground.

Once these experiments have been successful

it becomes of immediate interest to study the

spectrum and the distribution of the spectrum

over the disk under disturbed solar conditions

and during an intense flare. The solid propellant rocket and sea recovery are almost a

prime requisite for this purpose because they

make it possible to fire at almost the instant a

flare takes place. Though this may also be

possible with Aerobee-Hi, it will be difficult

indeed, owing to the safety requirements
42 SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS VOL.1

associated with launching from ground and

because of the great complications surrounding

the use of an acid-aniline fuel.

A first approach to the question of the change

of ultraviolet and X-ray emission during a

flare is planned by Friedman and his coworkers.

They will instrument a number of "Rockoons."

These small rockets are carried aloft at sea with

balloons, and launched from an altitude of

80,000 feet; they reach approximately 80 km,

which is sufficient for studies of Lyman alpha

and hard X-rays. When they are launched at

sea they can be watched during the entire day

and fired at the moment a flare appears. Each

rocket will carry a photon counter sensitive to

Lyman alpha only, another to detect the

possible presence of 1-2 A X-rays, and a third

device to search for a burst of 7-ray Bremstrahlung that may be produced in a shock

wave in the solar atmosphere associated with

the flare. It is also possible to have several

Rockoons ready to launch in sequence, a few

minutes apart, to study the time history of

the radiation from the flare. Such experiments can be done far better from a simple

solid propellant rocket, since this would carry

the instruments much higher than the Rockoon,

to altitudes where soft X-rays and regions of

the ultraviolet other than Lyman alpha can

also be studied. Furthermore, if no flare

should appear during the day the rockets would

not be launched, and so would not be lost, as

are Rockoons.

Still another interesting experiment would be

to study the extreme ultraviolet and X-ray solar

spectrum during a solar eclipse. It is not out

of the question to do this spectrographically

with solid propellant rockets launched at sea. A

further development would be required to enable

a pointing control to track accurately in spite of

the enormous change in solar light. It would

also be quite possible to photograph the flash

spectrum, and the results would be most important for deducing the correct model of the

chromosphere.

Friedman has proposed studying the sun

with photon counters and Rockoons during an

eclipse. Here one would hope to determine the

distribution of intensity in the corona of radiation in several bands of the X-ray region, and

the distribution of Lyman alpha with height in

the chromosphere and corona. This experiment would be much better performed with a

solid propellant rocket, since the balloon may

drift out of the eclipse path prior to totality.

The stars.—It would, of course, be of tremendous interest to use rockets to study the

short-wavelength spectrum of the stars. Though

their faintness in the visible makes the problem

appear extremely difficult, certain of the O and

B stars are so very hot that their extreme ultraviolet and X-ray spectrum may well be sufficiently intense. Exploratory investigations can

be accomplished with photon counters, and here

again Friedman has planned several experiments, both by day and by night. For the

first experiments the sky will be scanned, as the

rocket rolls and yaws, by a Lyman alpha photon

counter of high sensitivity and also by certain

other ultraviolet and X-ray counters. In later

experiments the counter will be directed with a

pointing control. By day the situation is complicated by the expected presence of scattered

Lyman alpha radiation from the sun, which in

itself is worth studying. At night, however,

Lyman alpha and X-rays may be expected from

stars, nebulae, and other galaxies, and there is

even the possibility of detecting these radiations

in the zodiacal light. The relation between the

F-corona and the zodiacal light could be studied

by visible or ultraviolet light photometry of the

sky from a rocket at altitudes where the Rayleigh sky background is very dim.

Many experiments can be imagined if an

accurate stellar pointing control were developed.

For example, A. Boggess and J. E. Milligan have

suggested photographing the stars in the ultraviolet with a Schmidt camera and objective

prism. The energy distribution of ultraviolet

spectra, worked out from these spectra, would

be of great interest in connection with stellar

atmosphere studies. It might also be possible

to extend knowledge of the wavelength dependence of interstellar reddening into the ultraviolet. Ultraviolet spectroscopy of the planets

might well give information concerning possible

absorbing atmospheric constituents. Perhaps

most exciting of all, however, is the possibility

that the spectra of many of the stars could be

photographed in the vacuum ultraviolet.

Particles.—There are other fields of astronomical research which can be pursued by means of
NEW HORIZONS IN ASTRONOMY 43

rockets. One of these is the study of micrometeorites. These tiny particles are thought to

bombard the outer atmosphere in such abundance that it may be possible to detect acoustically their impacts on a rocket, or even to

collect a sample and parachute it to earth for

study. Experiments of this sort are being

conducted by both the Navy and the Air Force.

Another field is the analysis of the primary

cosmic radiation. This is difficult at low altitudes because of the modifications and showers

that take place due to the stopping effect of the

air. At rocket peak altitudes, however, it

should be possible to study the real primary

radiation.

Upper atmosphere physics.—Finally, there is

the large field of research on the physics of the

earth's upper atmosphere, a field claimed alike

by astronomy, geophysics, and meteorology.

Rocket work has already solved the problem of

the origin of the E-layer, and when higher

flying rockets are available the radiation responsible for the F-layers will certainly be discovered. Direct measurements of the composition of the upper atmosphere have been made,

and the density of ions of various sorts has been

measured with mass spectrometers. The electron density through the E and Fi layer has

already been measured from rockets by a radio

propagation method by Seddon (1954) and

coworkers (Seddon et al., 1954). This work

should be carried on at many points over the

earth and under varying ionospheric conditions. The pressure and density of the atmosphere have been measured from rockets and

the work should be continued at many places

over the earth and at different seasons. It

may be noted that the density measurements

of Friedman (1955), made by studying the

absorption of X-rays by the atmosphere, do not

agree with measurements made close to the

rocket with pressure gages. This interesting

conflict of results must be resolved.

Finally, it is hoped that the origin of aurorae

can be explained by means of rockets. It is

expected that during the International Geophysical Year rockets can be launched within

the auroral zone and flown into actual aurorae.

The nature of the incoming particles will be

studied. Emission of Lyman alpha will also be

looked for, since it may be expected from recombining protons and electrons.

Outside the auroral zone spectacular displays

are infrequent, but there is still present the night

airglow, an emission of light from the upper atmosphere. The origin of the various emissions

in the airglow and their possible connection with

the aurora are not well understood; and opinions

differ concerning the altitudes at which the

various spectral lines are emitted. Experiments from rockets designed to study directly

these emissions are underway; they should be

continued until the airglow is well understood.

For all these experiments the development of

the solid fuel propellant rocket is most important. Since many experiments will be required,

at special times and locations, a cheap and

simple rocket is essential.

Satellite

A satellite vehicle is simply a special type of

rocket that is projected at a velocity proper to

place it in a stationary orbit. It is especially

suited to experiments that require a long observing time. In the foreseeable future, however, the space and weight limitations on instrumentation are very severe, and there can be no

thought of recovery.

There have been proposed several important

experiments that cannot be done with ordinary

rockets and may be feasible from a satellite.

Perhaps the simplest experiment is the study

of solar ultraviolet and X-ray radiation over a

long period of time in order to determine the

type of intensity variations that may occur.

Such variations are considered to be a cause of

sudden changes in the state of the ionosphere

and radio transmission, and may also have an

important effect on weather. Photon counters

would be used.

Friedman and coworkers have proposed a

search of the sky for radiation hot spots by

means of detectors in a satellite. The Lyman

alpha line would be the radiation chosen for

the first experiments because of its probable

high intensity and possible correlation with

known regions of high concentration of

hydrogen.

The earth's magnetic field could very well

be studied from a satellite. Since the satellite's
44 SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS VOL.]

orbit would cover a large fraction of the earth

during a day, it would be possible to study the

variation of the earth's magnetic field with

latitude and longitude, and, if the orbit were

elliptical, also with altitude. Experiments of

this sort are important in connection with ttfe

Chapman-St0rmer ring current. An actual

experiment from a satellite to determine the

magnitude of the ring current has been proposed

by J. P. Heppner.

The satellite would, of course, be ideally

suited to the study of micrometeorites so small

that they do not penetrate far into the earth's

atmosphere but which probably ionize the

upper atmosphere in process of being stopped.

Even more important, however, may be the

study of primary cosmic rays of energies so

low that they do not penetrate far into the

atmosphere. Particles of this sort are believed

to exist and, in fact, many of them may come

from the sun. The satellite also offers the

possibility of studying the distribution in

space from which cosmic radiation arrives, and

perhaps of providing clews to the origin of this

mysterious radiation.

References

BEHRING, W. E.; JACKSON, J. M.; MILLER, S. C, JB.;

AND RENSE, W. A.

1954. Journ. Opt. Soc. America, vol. 44, p. 229.

BOYD, R. L., AND SEATON, M. J., EDS.

1954. Rocket exploration of the upper atmosphere. Interscience, New York.

BYRAM, E. T., CHUBB, T. A., AND FRIEDMAN, H.

1955. Phys. Rev., vol. 98, p. 1594.

FRIEDMAN, H.

1955. Ann. de Geophys., vol. 11, p. 174.

FRIEDMAN, H., AND CHUBB, T. A.

1955. In Report of the 1954 Cambridge Conference, pp. 58-62. The Physical Society, London.

HAVENS, R. J., FRIEDMAN, H., AND HULBURT, E. 0.

1955. In Report of the 1954 Cambridge Conference, pp. 58-62. The Physical Society, London.

JOHNSON, F. S.; MALITSON, H. H.; PURCELL, J. D.; AND

TOUSEY, R.

1955. Astron. Journ., vol. 60, p. 165.

JURSA, A. S., LEBLANC, F. J., AND TANAKA, Y.

1955. Journ. Opt. Soc. America, vol. 45, p. 1085.

RENSE, W. A.

1953. Phys. Rev., vol. 91, p. 299.

SEDDON, J. C.

1954. Journ. Geophys. Res., vol. 59, p. 463.

SEDDON, J. C, PICKAR, A. D., AND JACKSON, J. E.

1954. Journ. Geophys. Res., vol. 59, p. 513.

STACEY, D. S.; STITH, G. A.; NIDEY, R. A.; AND PIETENPOL, W. B.

1954. Electronics, vol. 27, p. 149.

TOUSEY, R.

1953a. In Kuiper, ed., The sun, p. 658. University of Chicago Press.

1953b. Journ. Opt. Soc. America, vol. 43, p. 245.

WILSON, N. L.; TOUSEY, R.; PURCELL, J. D.; JOHNSON,

F. S.; AND MOORE, C. E.

1954. Astrophys. Journ., vol. 119, p. 590.
Computing Machines in Astronomy

By Paul Hergetl

The impact of computing machines upon

scientific research in recent years is not fully

appreciated outside of a small group of actual

machine users. When an electronic computation center is installed in a scientific or technical

institution, almost invariably a workload

quickly develops that requires more than the

available running time on the computer. This

is a "snowballing" situation in which the demonstration of what can be done on one problem

soon leads to the exploration of similar possibilities in related problems.

The next effect is that nearly all nondefense

and pure research projects are crowded out, for

lack of priority. The National Science Foundation has recognized this situation and called

an ad hoc committee in February 1955 to consider the need for computing facilities in

university research centers in general. Astronomy was represented by Schwarzschild and

Herget. More recently other groups have

recognized the need to stimulate the development of computation centers in universities.

The need for electronic computing machines

in classical astronomy is mentioned in another

paper in this series, "The Minor Planets."

As additional examples, we may note that

Brouwer, Clemence, and Eckert, using these

machines, have completely recomputed the

motions of the outer planets and Herget has

prepared a uniform system of coordinates for

Earth and Venus. Clemence has recomputed

a new general theory for Mars which is nearing

completion. In studying the secular variations

of the elements of the major planets, it is now

feasible to compute second order effects simply

by recomputing the coefficients under the

appropriately varying conditions at intervals

of, say, 1,000 years. An extensive computational program is also a significant part of the

1

 Department of Astronomy, University of Cincinnati, Cincinnati,

Ohio.

meteor studies currently being conducted by

Whipple.

In astrophysics, one large group of problems

concerns stellar models. The equations of

state and equilibrium conditions for stellar

interiors form a system of ordinary differential

equations in one independent variable, the

radius. Similar mathematical but different

physical conditions exist in stellar atmospheres.

Recently Henyey has elaborated on the problem of stellar interiors by taking into account

the evolutionary changes within the star.

Again, if one deals with pulsating stars, not

only the stellar radius but also the time must

be used as an independent variable. Then

there are such specialized problems as the search

for a mechanism that can explain the occurrence

of spicules on the sun.

Another new field of computational problems

is now arising from considerations of magnetohydrodynamics. These are even more difficult

because they involve partial differential equations. The whole field of "wave equations"

is of this same type. Much work has been done

by L. C. Green, L. H. Thomas, and others, but

it is always necessary to restrict the equations

somewhat because even the largest computers

in existence are still inadequate to deal with

this problem in its entirety. This problem is

of interest not only for the light astronomical

atoms but more or less for the whole periodic

table up to the fissionable and heaviest elements.

Examples of problems in stellar atmospheres

may be found in the Proceedings of the National

Science Foundation Conference held at Indiana

University in October 1954. Other typical

problems are given in the report of the Joint

Meeting of the American Astronomical Society

and the Association for Computing Machinery

held at Ann Arbor, Mich., June 1954. One

interesting problem was the computation of the

history of a globular cluster (in two dimensions)

48
46 SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS VOL.1

in which the individual stellar trajectories were

carried forward simultaneously until, eventually, two out of a hundred stars evaporated

from the cluster.

Optical design, lens computations, and raytracing have been revolutionized by the development of high-speed computers. Thus the

cost of experimenting with a novel design is

greatly reduced, since all the effects can now

be so readily computed. Highly efficient methods have been developed, especially by Baker,

but there is still room for further study as the

capabilities of newer machines are improved.

The design of radio antennas for special

purposes is another new computational problem

in astronomy.

It is significant that, in nearly every one of

these astronomical problems, the researcher

has been able to develop his own facility in

the use of an electronic computer as a research

tool even though he has not been primarily a

computer. The use of computing machines in

astronomy will undoubtedly grow exponentially

as each successful solution breeds more problems, and each new experience inspires more

ideas.
Turbulence

By S. Chandrasekhar1

An aspect of theoretical investigations in astronomy and geophysics that is not often recognized is that matter and motions, in the large,

may exhibit patterns of behavior that one might

never suspect if one restricted oneself to matter

and motions in the small. While several examples can be cited to illustrate this, the most

striking is that which requires all motions in

systems with large linear dimensions to be turbulent; for turbulence will be the natural state

of motions if the Reynolds number is sufficiently large. And since the Reynolds number contains the linear dimension of the system as a

factor, it is clear that hydrodynamical motions

which occur in astrophysics and geophysics must

necessarily be turbulent. Indeed, "turbulence"

is the refuge under which astronomers most

commonly seek protection whenever they do not

understand a particular phenomenon. At present, however, our understanding of the nature

« Yerkes Observatory, Williams Bay, Wls.

of turbulence is no greater than was our understanding of the dynamical theory of gases, before Maxwell and Boltzmann.

Since a deductive fundamental theory of

turbulence is a prerequisite for a rational comprehension of many astrophysical and geophysical phenomena, an attempt to follow current developments in the theory of turbulence

should be an integral part of astronomical education. The need is the greater in view of the

increasing importance of magnetism in astronomical thinking; we have only to recall in this

connection the problems of interstellar polarization and magnetic fields, the synchrotron

radiation from the Crab Nebula, tie galactic

corona, and the still unsolved problem of the

earth's magnetic field.

Hydrodynamics and hydromagnetics are

bound to play central roles in future astronomical developments and the part of turbulence is

only one aspect of this.

47

Astroballistics

By Richard N. Thomas

 1

Over the postwar years a small group of

astronomers have found themselves acting as

midwives at the birth of a new subject, astroballistics, conceived on the common borders

of astronomy, aerodynamics, ballistics, and

chemical reaction-kinetics. Its multiple origin

provides an excellent example of the preoccupation with crossfield research in contemporary

astronomy. Its evolution exhibits the trend,

beginning with spectroscopy, which is changing

astronomy from a wholly observational and

theoretical science to one demanding extensive

laboratory investigations. Its future depends

on our success in establishing this crossfield

work; that is, on our continuing to produce a

group of investigators who feel at home in at

least the overlapping parts of the four areas—

astronomy, aerodynamics, ballistics, and chemical reaction-kinetics.

Origin and scope

The motion of planets and planetoids was for

years a conventional part of study in astronomy.

In other papers of this group appear descriptions

of some of the exciting new data and theories

which have deflected our interest in these

objects from their kinematic to their physical

properties. We recognize, however, that we can

hardly claim a satisfactory knowledge of our

planetary system until we can also specify

the properties of those objects so small that

they cannot be observed individually. In

general, our information on these objects is

wholly statistical, coming from studies of their

collective scattering effect on sunlight, action

as a resisting medium, etc. One notable exception is the study of meteors; for these, the earth

serves as a sampling box, sweeping up the

interplanetary material, with the earth's atmosphere as a fluorescing medium, responding

to the entrance of each meteor. If we com1

 Harvard College Observatory, Cambridge, Mass.

pletely understood all the processes that occur

when a solid object moves through a gaseous

atmosphere, we would have a fine method for

inferring the corresponding properties of each

meteor observed. We do not yet have such

an understanding, however, and the science of

astroballistics has arisen from our attempts to

remedy this lack. Our present knowledge of

meteors, as a part of the interplanetary medium,

is discussed elsewhere in these papers. Here

we consider the meteor only as it represents a

solid object moving through a gaseous atmosphere, under one variety of conditions.

When a solid moves through a gas, several

things may happen. The solid loses momentum to the gas and is decelerated. The amount

of deceleration depends critically upon the

mass, shape, and velocity of the body, and upon

the composition and thermodynamic state of

the gas. The solid loses energy as it decelerates, and the energy contributes heat to both

body and gas. The partition of heat between

the two depends critically upon the velocity,

shape, composition, surface structure, and

thermodynamic state of the body, and upon the

composition and thermodynamic state of the

gas. The heat transfer to the solid may result

in thermal radiation and ablation (either by

melting or by vaporization) from the surface, as

well as in a simple temperature rise throughout

the interior. The heat transfer to the gas may

result in excitation and ionization, and thus in

radiation, if the temperature rises sufficiently.

Moreover, the material ablated from the solid

may interact chemically with the gaseous atmosphere, serving either as a heat source or

sink, depending upon the reaction.

The astronomer can make at least three types

of measures on a meteor moving through the

atmosphere: (1) He can measure directly time

and position, and so obtain velocity and

deceleration; (2) he can measure the total radi48
50 SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS VOL.1

ation, and possibly its spectral distribution; and

(3) by radar, he can infer the ionization produced by the meteor. Ideally, he would like to

have available a set of tables which he could

enter with these data, and so read off the physical characteristics of a single body which alone

could have produced such a disturbance. At

worst, he might expect to obtain a range of such

bodies, and this range would represent his uncertainty of the properties of the particular meteor observed. These properties he could compare with those of rare meteors that were large

enough and slow enough for the bodies to survive passage through the atmosphere, and be

recovered as meteorites.

Unfortunately, when we attempt to find or

even to prepare such a set of tables, we find a

disheartening situation. Meteors impact on the

earth's outer atmosphere at speeds ranging between 10 and 75 km sec"1

. The ballistician and

aerodynamicist, to whom we might expect to

turn for information, have just since the war

become really conscious of the velocity range

that is of interest to the meteor astronomer.

They, in turn, have looked to the astronomer,

hoping for insight into just those problems we

pose to them. Only during the last war did the

aerodynamicist develop a keen interest in speeds

as high as that of sound, ~% km sec"1. While

the ballistician customarily dealt with speeds

several times that of sound, his knowledge was

primarily empirical, and his empirical investigations were limited by bis accelerating devices.

The experimental investigations reported in the

unclassified literature do not deal with objects

of known geometry whose speed exceeds ~ 2 km

sec"1. At these speeds, interest lies in airresistance; aerodynamic heating of the body is

negligible. By analogue devices such as the

wind tunnel, deceleration properties of bodies of

known shape at equivalent airspeeds of ~ 3 km

sec"1 have been studied. Theoretical investigations have made rapid progress in interpreting

and correlating these empirical studies. On the

whole, we may say that the deceleration problem is in a most satisfactory state, not only in

the range covered by the experiments but, for

purposes of extrapolation, considerably beyond

that range. But deceleration measures alone do

not suffice to permit an inference of the physical

properties of the meteor. For that, one must

turn to luminosity and ionization studies; and

here, the cupboard is essentially bare.

Ballistic and aerodynamic studies of objects

moving at such high speeds that ablation occurs

are virtually nonexistent, at least so far as we

can infer from the unclassified literature. A

few simple free-flight experiments have been

made, with velocities reaching ~ 6 km sec"1 at

maximum. We find an equally small number

of wind-tunnel studies, involving always substances requiring but low heat transfer for

ablation. From a theoretical standpoint, the

situation is equally bad. The study of the

aerodynamic flow pattern about a body that is

losing mass at its leading surface has hardly

been discussed. Indeed, the most satisfactory

theoretical work on any aspect of the field

probably lies in a few studies, made in collaboration by astronomers and chemists, relating

to ablation in the meteoric regime—high

velocity and very low gas density.

When we consider the data available on

excitation and ionization caused by the ablated

material ejected into the gaseous atmosphere,

we find the bleakest situation of all. In the late

1920's and early 1930's theoretical atomic

physics was striking out into the areas opened

by the new quantum mechanics, and experimental physics was entering those areas opened

by improved instrumentation. Problems of

radiative transition probabilities, while not

completely solved, were at least understood.

On the other hand, problems of collisional

excitation were vague, expecially those involving exchange reactions, and particularly

atom-atom inelastic collisions. A number of

preliminary experimental investigations were

reported in the energy range of interest to the

astronomer, < 1 kev. Unfortunately, from our

present standpoint, the early 1930's were also

the era of the crucial experiments in the then

infant science of nuclear physics. The focus of

interest rapidly shifted, and concern with this

area of low-energy atom-atom collisions diminished. An extensive base of data upon

which to draw never developed, and it is only

in recent years, among relatively small groups

of physicists and chemists, that interest in these

problems for their own sake has reawakened.

The astronomer's interest arises largely from

his meteor studies, and occasionally from
NEW HORIZONS IN ASTRONOMY 51

phenomena in stellar atmospheres where, however, his chief concern lies in electron-atom

inelastic collisions. As the instrumental range

increases, particularly in studies of very strong

Shockwaves, the aerodynamicist also is increasingly preoccupied with the details of these

atomic collision processes, excitation, ionization, and thus radiation.

To categorize and collate the variety of problems just described, the science of astroballistics developed. Its province includes the

study of those phenomena arising out of the

motion of a solid through a gas at such speeds

that ablation occurs. Interest in and data on

various of these phenomena arises in astronomy,

aerodynamics, ballistics, and chemical physics.

Investigations that combine the points of view

of the several sciences involved will have a

better perspective and are more likely to solve

the various problems than a study based on

any individual science alone.

Evolution of the science

The field of astroballistics, as such, developed

when astronomers recognized that it would

probably be necessary to set up a systematic

program that linked astronomical observations

of meteors with laboratory studies in ballistics

and aerodynamics. The immediate postwar

period found ballisticians and aerodynamicists

still preoccupied mainly with measures of the

system of aerodynamic forces on a missile or

aircraft and, to a lesser extent, with problems

of strong Shockwaves. A number of advances

in experimental techniques, however, made it

possible to begin, at least, to bridge the gaps

between meteor observations and the data

and theory of ballistics and aerodynamics.

A certain amount of development of new

launching devices, and considerable improvement of the old, plus the use of low meltingpoint materials, together have made it possible

to simulate ablation from meteors. Newly

developed free-flight firing ranges permit studies

under varying conditions of the atmosphere

through which the solid moves. Advances

in high-speed photography allow us to make

inferences of the spectral distribution of luminosity along the axis and trail of an artificial

meteor, which we have been able to accelerate

up to speeds of 6 km sec"1. Microwave equipment facilitates studies of the ionization

caused by the same type of artificial meteor.

Theoretical studies already begun, based on

the insight developed in these first investigations, will help guide subsequent research.

Thus, we are laying a solid foundation from

which to study the problems described earlier

as largely unsolved. We should, however,

stress one aspect of all these investigations:

the results of these first attempts have been

invariably crude and preliminary, and progress

in extending them has been dishearteningly

slow. We will return to this point. At

present, the one satisfying feature the results

have in common is the simple fact that they

exist.

These first experimental investigations had

one very stimulating outcome. The first

spectral observations, though crude, demonstrated that it is possible to observe details

of the chemical reactions involving ablated

material and atmospheric atoms. Meteor observations have never provided such data,

primarily because of the unfavorable conditions

of (1) high velocity, which inhibits the exothermic chemical reactions expected when a

metallic spray hits the atmosphere; and (2)

low atmospheric density, which permits considerable diffusion before conditions are such

that the reaction may occur. In the laboratory,

such inhibition and diffusion effects may now

be studied. Thus, a very fertile method for

studying chemical reaction kinetics appears.

Even more, all these occur in a high-speed

stream of gas, so that we have an approach to

an aerodynamic flow problem of exciting

possibilities. For the astronomer, who is just

now acquiring a deep interest in gas-dynamical

problems in stellar atmospheres, such observations are invaluable in providing a physical

insight otherwise inaccessible. In the laboratory, it is hard to provide gas streams

having sufficient internal energy that atomic

and molecular excitation change appreciably

and observably. Consequently, such data as

the above become highly desirable. The remark made above, however, should be repeated; these studies are chiefly preliminary.

The field has just been opened.
52 SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS VOL.1

Needs and future of astroballistics

The gross problems of astroballistics may be

clearly delineated. A very great deal of

experimental work must take place, both to

provide the empirical constants required in any

theoretical or interpretive work and to establish the basic physical feeling for the problems

without which the theory cannot proceed. At

present, facilities for such work are largely

nonexistent or inadequate in comparison with

our needs. Consequently, the field can develop

only slowly.

For a long time the chief characteristic of

American astronomy has been the size, quality,

and quantity of its telescopes. Astrophysics in

America has grown largely through the output

of observational data. Several spectroscopy

laboratories are contributing invaluable work

on wavelengths and transition probabilities.

Aside from these, however, we have no corresponding astronomical laboratories devoted

to the study of astroballistics, of conditions

underlying particle growth in the interstellar

medium, and of magnetohydrodynamic properties of plasmas and the like. It is of course

true that some phase of each of these subjects

may be under investigation at a dozen places

throughout the country. But in how many

of these places is there an astronomer working

as a member of the research group, both contributing and gaining information ? How much

of the information gained is accessible in the

open literature? We should recognize two

problems here. First, we have need for a

national astrophysical laboratory that can operate with considerable latitude to investigate

problems such as those in the field of astroballistics. Its staff should comprise astronomers

as well as scientists from "related" fields, and it

should include both permanent members and

those on a temporary basis. Second, we must

urge our graduate schools to help produce the

staff for such an institution by emphasizing in

their program the experimental aspects of

astronomy as much as the observational.

Astroballistics is hardly unique in requiring

these supports; it is, after all, but a small field

in astronomy. The needs of astroballistics,

however, are a prototype of those existing in

contemporary astronomy in general, and the

future progress of the field may serve as a

significant example.
Electrodynamics of Fluids and Plasmas

ByMaxKrook1

Sufficient evidence has accumulated to establish the widespread occurrence of largescale electrodynamic phenomena in astrophysical contexts. To understand such phenomena and to interpret them quantitatively

we must investigate the origin, nature, and

behavior of cosmic electromagnetic fields and

their interaction with cosmic matter. These

tasks naturally involve considerations of fluid

electrodynamics and, from a more general and

fundamental point of view, considerations of

the kinetic theory of plasmas.2 Since these

latter branches of physics are still in a very

rudimentary stage of development, there is

considerable scope for research not only in

astrophysical applications but also in the

fundamental physical disciplines themselves.

An appreciable fraction of cosmic material

consists of ionized, and hence electrically conducting, gas. The dynamical behavior of such

conducting fluids is determined in part by forces

exerted by the electromagnetic field on electrical

charge and current distributions in the fluid.

But equally, the dynamical behavior of the

electromagnetic field is determined in part by

the fluid motions. In general therefore, conducting fluid and electromagnetic field have to

be treated simultaneously and on an equal

footing as interacting components of a system.

The study of systems of this kind is termed

fluid electrodynamics (or magnetohydrodynamics

or hydromagnetics) when the fluid is treated as a

continuum, and is termed kinetic theory of

plasmas when the fluid is treated microscopically.

Our aim here is to emphasize unsolved problems and to discuss possible future trends. No

attempt, therefore, is made to review the

1

 Harvard College Observatory, Cambridge, Mass.

1

 In this article the term "plasma" will be used in a general sense to

mean any partially or wholly ionized gas.

381014—56 5

literature in any comprehensive way. In a

discussion of this kind, it is not possible to do

justice to the individual contributions of the

many authors in these fields, but some published

reviews which include extensive bibliographies

are listed at the end of this article.

Fundamental equations

In the kinetic theory of plasmas, the state of a

system is specified by the following:

(1) State-variables for the electromagnetic field.

These are the usual parameters of the Maxwell-Lorentz

theory; i. e., electric intensity E, magnetic intensity H,

charge density p«, current density J.

(2) State-variables for the matter. These are

parameters which give a statistical description of the

microscopic state of the plasma.

With any dynamical process in a plasma is

associated some characteristic length L, and

some characteristic time T. If L and r are

both "sufficiently large," the plasma may be

treated as a continuum whose dynamic and

electrodynamic properties are specified by a

set of material constants; e. g., dielectric constant K, magnetic susceptibility /x, kinematic

viscosity v, electrical conductivity <r (which

may have directional properties), etc. In this

macroscopic theory, applicable also to conducting liquids, the state of the fluid is specified

by the usual variables of fluid dynamics, i. e.,

pressure p, density p, temperature T and flow

velocity©.

The fundamental equations of the continuum

theory are made up of the Maxwell electromagnetic equations together with the equations

of fluid dynamics, but with the following

modifications (which represent coupling between

electromagnetic field and fluid motion):

(a) The momentum and energy equations of the

fluid include terms which involve electromagnetic

53
54 SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS VOL.1

variables; e. g., a form for the momentum equation is:

where F=external force density, and c=velocity of

light.

(b) The constitutive equation for the current

density is altered, e. g., to the form:

On account of the complexity of this system

of equations, it is generally extremely difficult

to solve dynamical initial- and boundary-value

problems of fluid electrodynamics, and especially

so in cases of astrophysical interest. For this

reason, effort has been concentrated on investigating simple and idealized types of problems

and on exploring the consequences of various

kinds of assumption. It is hoped that insight

into the electrodynamic behavior of fluids gained

in this way may shed some light on the real

dynamical problems and eventually lead to

their solution.

Some general principles of the subject may

be obtained by analyzing the mathematical

structure of the fundamental equations. Dimensional considerations, for example, show

that significant mutual interaction of fluid

motion and magnetic fields can occur only

when the "scale" of the system is large enough.

For systems of ordinary laboratory size we

would generally require high magnetic intensity,

or high electrical conductivity, or low density,

or some combination of these factors. This

explains why magnetohydrodynamic behavior

appears more naturally in astrophysics (where

linear dimensions of systems are large) than in

terrestrial physics.

The development of fluid electrodynamics

would be advanced immeasurably by the

elaboration of techniques for performing controlled experiments on a laboratory scale and

over a wide range of behavior. The experiments so far carried out with liquid mercury

and liquid sodium barely touch the fringe of

the problem; and similar remarks apply also

to plasma kinetics. Although certain aspects

of plasma physics have been studied extensively

in laboratory experiments, e. g., with electrical

discharge tubes, many dynamical aspects of

particular interest in astrophysics remain

practically unexplored.

Ionospheric studies are also capable of

yielding valuable information about plasma

dynamics. And conversely, plasma dynamics

can play an important part in interpreting

ionospheric behavior. The structure of the

earth's upper atmosphere can be studied

extensively and accurately by various

techniques. Quantitative predictions of the

theory can therefore be subjected to more

stringent checks in the case of the ionosphere

than in most astrophysical cases.

Cosmic magnetic fields

In 1908, Hale discovered the strong magnetic

fields of sunspots (ii7~2,000 gauss) by utilizing

the Zeeman splitting of spectral lines. Great

improvement in this technique, particularly

by H. D. and H. W. Babcock, has resulted in

the detection of fluctuating fields in the solar

atmosphere (H~ l to 20 gauss), and of strong

variable fields in many stars (H~6,000 gauss).

The polarization in the light of distant stars

and its correlation with reddening have been

attributed to orienting effects of interstellar

magnetic fields on elongated dust particles

(H~10~5 gauss). A circumstance which is

probably significant for the theoretical interpretation is that astrophysical magnetic fields

generally vary with time.

The coupling of material motions and electromagnetic field implies that, in a conducting

fluid, mechanical energy and electromagnetic

energy are convertible one into the other.

Since relative motions are commonplace in

astrophysical plasmas, we can understand in

a rough qualitative way why magnetic fields

might exist there. (In a conducting fluid, magnetic field energy is generally much greater than

electrical field energy, except possibly under

transient conditions.) Quantitatively, and

even from a detailed qualitative point of view,

however, we are still far from having a satisfactory theory of astrophysical magnetic fields.

Order of magnitude estimates are often based

on a so-called equipartition theorem, which

asserts that the average density of kinetic

energy \po2 in a system equals the average

density of magnetic energy fiH2/8ir. For the
NEW HORIZONS IN ASTRONOMY 55

interstellar gas, this would imply fields of the

order 5X10"6 gauss. The theorem has so far

been verified only in a small number of comparatively trivial cases, and its real range of

validity is not known. Even for systems in

which the theorem could be true, the equipartition state may not be possible because of

very long relaxation times; this question can be

decided only by examining the mechanisms

which operate to convert mechanical energy

into magnetic energy and vice versa.

Our knowledge of such mechanisms is still

severely limited. On general grounds it appears that an efficient process for amplifying an

initially small magnetic field might be provided

by the interaction of turbulent fluid motions

with the field. Very little progress has yet

been made in the study of magnetohydrodynamic turbulence, especially in astrophysically

interesting cases of compressible fluids. Much

of the difficulty of this subject is due to the fact

that the sort of questions in which we are

interested cannot reasonably be answered by

combining a statistical theory of turbulence

with a dynamical theory of the electromagnetic

field. When a dynamical theory of turbulence

is eventually developed, great strides will be

possible in the theory of magnetohydrodynamic

turbulence.

On account of many uncertainties in our

knowledge of the magnetic properties of interstellar dust, our knowledge of galactic magnetic

fields is correspondingly very inexact. It is

therefore important to look for further observable consequences of galactic fields which might

lead to more exact and detailed information on

their structure and behavior.

One promising line of approach appears to be

through problems of generation and propagation of primary cosmic ray particles in the

galaxy, although it is still too early to anticipate

where, or how far, this approach will lead.

These problems present a wide and attractive

field for future research.

Preliminary calculations with the Fermi mechanism have yielded some results which are consistent with the observational data and some

results which are not. In this model, which is

capable of some elaboration, particles of sufficiently large initial energy increase that energy

by successive encounters with magnetic fields

"frozen" into interstellar gas clouds. The lowenergy cosmic ray particles are supposed to be

injected into the galactic accelerator by stars,

and this is not inconsistent with recent evidence

which indicates that the sun itself may be

ejecting low-energy cosmic rays.

It has recently been found that optical, continuous radiation from the Crab Nebula is

strongly polarized. This has been interpreted

as synchrotron radiation from electrons in

magnetic fields of order 10~* gauss. An intensive study of this and similar nebulae in both

optical and radio frequencies may well prove

invaluable in our understanding of mechanisms

for generating cosmic rays.

Magnetic fields and other electrodynamic

phenomena can play a significant role in

problems of structure, behavior, and evolution

of astrophysical systems. There is considerable

scope for research in problems of this kind.

The structure of a star may, for example, be

modified appreciably in regions where magnetic

pressure is comparable with gas pressure; also,

magnetic fields may have profound effects on

the stability of configurations. Further, a

general galactic origin of cosmic rays, as in

Fermi's model, implies a heavy drain on the

mechanical energy of gas clouds; this would

have a substantial effect on galactic evolution.

Kinematic models

Since solutions of proper dynamical problems

are difficult to obtain, kinematic models can

play an important part in the development of

the theory. In the kinematic approach we

assume that the motion or configuration is of

a particular type and then examine the consequences with the fundamental equations. The

dynamo mechanisms proposed by Elsasser and

Bullard to account for the permanent field of

the earth are of this kind, as are models of

torsional oscillations of the sun in its own

magnetic field. The most fruitful line of attack

on questions of the origin and behavior of sunspots appears, at present, to be through investigations of a variety of kinematic models.

In general, kinematic models cannot properly

be accepted as solutions of real dynamical

problems until their consistency has been

established by stability investigations. Such

stability problems are usually rather difficult,
56 SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS VOL.!

so the validity of most of the existing kinematic

solutions has remained untested.

Linearized equations

The nonlinearity of the basic equations is a

major source of difficulty in the theory. A

linearized form of the equations may be used

for treating motions of small amplitude; e. g.,

propagation of magnetohydrodynamic waves.

Limited but valuable information can be

obtained from the discussion of such problems,

especially for systems with boundaries and with

gradients of the physical parameters; e. g.,

density.

One of the more important applications of

linearized theory is in the investigation of

stability of systems against small disturbances.

The stability of a number of simple systems in

external magnetic fields has been examined

in this way, but very little is known about

stability of systems for which the magnetic

field is wholly internal (cf. kinematic models).

Kinetic theory of plasmas

The material constants <r, v, etc., that appear

in the continuum theory are not determined

by that theory but have to be obtained by

experiment or from a different theory. A

kinetic theory of plasmas, on the other hand,

implicitly includes the phenomena of electrical

conduction, viscosity, etc., in its basic equations.

Such a theory is in fact invoked to calculate

<r, v, etc., for use in the macroscopic theory.

In addition, however, there are many cases of

interest, characterized inter alia by "sufficiently

small" values of L or T, for which the continuum theory is not valid and for which we

then have to use kinetic theory.

Several problems of a fundamental nature

already occur in the very task of formulating

basic equations for a kinetic theory of plasmas.

The Liouville equation, although useful for

certain formal investigations, is far too complex

to serve as a basis for attacking practical

problems. At the opposite extreme is the very

crude model of a plasma as consisting of intermingled positive, negative, and neutral fluids

with frictional coupling. Between these extremes there are many possible formalisms with

varying degrees of generality.

At present, the most fruitful compromise

between generality and mathematical tractability is the theory based on the use of singleparticle distribution functions fa(vl, 7, t).

(^=velocity of particle of type "a".) Particle

interactions have then to be represented in

idealized form—"close" interactions as binary

collisions, and "distant" interactions as average

internal electromagnetic forces. Static and

dynamic effects of particle correlations are

neglected in the first instance but allowed for

roughly by the introduction, post hoc, of the

Debye screening radius. The basic equations

of this theory consist of the Maxwell equations

together with a set of Boltzmann equations or

of Fokker-Planck equations; e. g., for particles

of type "a":

mc

where ea=charge, ma= mass of particle, ( ~ \

is a sum of Boltzmann collision integrals, and

E, H include contributions of internal fields.

Investigations of a fundamental character to

determine the range of validity of the theory

would be highly desirable.

In addition to the length L and frequency

u=— which characterize a particular dynamical

process, the plasma is characterized by a set of

intrinsic fundamental lengths and fundamental

frequencies. Typical frequencies associated

with particles of type "a" are the plasma

frequency <>>TM = {kxnae^lma]K, the gyro-frequency «i*)

=——> and the collision frequencies

with particles of type "j8", o#},; (na—particle

density). The modes of dynamical behavior of

a plasma depend intimately on the relative

magnitudes of the various characteristic lengths

and frequencies. These modes have not yet

been sorted out in a completely satisfactory

way for the general case.

The macroscopic theory can be derived from

the kinetic theory by the Chapman-Enskog

procedure in the limiting case of "high" density

(e. g., the <*£), » all other fundamental frequencies), and "small" gradients (e. g., L »

all mean free paths). To any type of problem

originally formulated in terms of the continuum
NEW HORIZONS IN ASTRONOMY 57

theory, there corresponds a more general formulation in terms of kinetic theory. The

general microscopic problem is of course more

complex than the special continuum problem.

Some simplification is introduced in another

limiting case, that of "low" density (neglect of

collisions), and "small" gradients (linearized

theory). For a gas of neutral molecules this

limit would correspond to linearized free-molecule flow; for a plasma, however, cooperative

behavior persists through the agency of internal

electromagnetic fields.

Linearized kinetic equations

The nonlinearity of the equations is again a

source of complication in the kinetic theory of

plasmas. The development of the theory as

reported in the literature has been confined

almost wholly to the linearized theory and, in

particular, to problems of small-amplitude

plasma oscillations. Such oscillations can be

studied comparatively easily in the two opposite

limiting cases of high density (collision

frequency dominant), and of low density

(collisions neglected); the intermediate range

between high and low density presents difficulties even though the equations have been

linearized.

For an infinite, uniform plasma in thermodynamic equilibrium, dispersion relations (i. e.,

relations between frequency « and wave number

k) have been obtained for the above two limiting cases. Qualitative and semiquantitative

results have been found for the intermediate

range by using approximate kinetic equations

which incorporate the physical properties of

the Boltzmann equation but are mathematically

more tractable. In the absence of external

fields, relatively undamped sound waves at

high density go over continuously into relatively

undamped, longitudinal (electron) oscillations

at low density, with heavy damping in the

intermediate region. The character of the

oscillations in the presence of external fields

has not yet been studied in complete detail.

In this case transverse and longitudinal waves

are in general coupled; magnetohydrodynamic

waves are the high-density analogue of lowdensity, free electromagnetic waves, etc.

Some attention has also been devoted to the

investigation of systems consisting of beams

of charged particles interacting with one another

and with a plasma. The examination of the

stability of such systems is involved in the

question of excitation of plasma oscillations

by such beams.

There is considerable scope for development

of plasma electrodynamics even within the

framework of the linearized theory. For the

astrophysical applications, we are particularly

concerned with boundary-value problems, with

the effects of gradients in the physical conditions (density, etc.), and with modes of exciting

plasma oscillations; e. g., by supplying mechanical energy. An important question, still

unresolved, is that of conditions under which

oscillations in a bounded plasma can provide

significant radiation in radio frequencies. Even

though the radio emission of objects like the

Crab Nebula may eventually be shown to consist predominantly of synchrotron radiation,

plasma oscillations probably will still be important in accounting for radio noise from other

sources; e. g., solar bursts, and for various

other electrodynamic phenomena.

In order to interpret observations of radio

frequency emission by the sun and other

sources, we require to know how electromagnetic

waves are propagated in plasmas. Such propagation problems are usually discussed in

terms of the magneto-ionic theory. They can

be treated more consistently and generally in

terms of the kinetic theory dealt with above,

although the two methods give concordant

results over a wide range in the small-amplitude

case.

Nonlinear phenomena

Since the fundamental equations are nonlinear,

plasmas are capable of types of dynamical

behavior not subsumed in the linearized

theory; e. g., the generation of discontinuities

from initially continuous motions. These

modes of behavior have not yet been investigated in any systematic way.

Many of the astrophysical phenomena to

which we would wish to apply the theory

(either microscopic or macroscopic) are so

violent as to make it extremely unlikely that

an adequate treatment of the phenomena would

be possible in terms of linearized theory.

Nonlinear plasma electrodynamics presents
58 SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS VOL.1

a whole new range of possibilities in our attack

on various astrophysical problems; e. g.,

generation of cosmic rays, generation of radio

noise, etc. On account of the intrinsic complexity of the subject, it will probably be

necessary to examine first a large number of

simple and idealized cases in order to gain

insight into nonlinear modes of behavior.

A rather limited range of nonlinear behavior

can be described in terms of oscillations of the

linearized theory with cross-modulation that is

determined by treatment of the nonlinear terms

as a perturbation. There are several interesting

problems of this kind concerned with the propagation of electromagnetic waves of "moderate"

amplitude through plasmas. Some of these

have a direct bearing on ionospheric propagation.

A typical nonlinear phenomenon, not reasonably representable in terms of cross-modulation,

consists in the generation and propagation of

shock waves. Some aspects of magnetohydrodynamic shocks have been discussed in the literature. To study the physical processes which

may occur in the presence of shock waves in a

plasma, we must examine the detailed microscopic structure of shock fronts and their stability. Such investigations must be based, of

course, on the full nonlinear equations and are

generally rather complicated.

Certain variations of the magnetic field of the

earth have been attributed by Chapman and

Ferraro to clouds or streams of ionized gas

ejected by the sun and impinging on the earth's

magnetic field. Such clouds are also supposed

to be responsible for various other observed

phenomena; e. g., aurorae. Very little is known

about the dynamical behavior of such clouds and

then* interaction with external fields. Still less

is known about the particular kinds of instability in the sun's atmosphere which result in

the ejecting of clouds and about the dynamics

of the ejection process.

Our discussion of problems in the electrodynamics of fluids and plasmas is by no means

exhaustive. Many aspects have been treated

only in cursory fashion; other aspects have not

even been mentioned. The development of the

subject has barely begun and the field is a wide

open one.

General references

ALFV£N, H.

1950. Cosmical electrodynamics. Oxford University Press.

BIKRMANN. L.

1953. In Beckerley, ed., Annual review of

nuclear science, vol. 2, p. 335. Annual

Reviews, Inc., Stanford, Calif.

CHAPMAN, S., AND COWLING, T. G.

1952. Mathematical theory of non-uniform gases.

Cambridge University Press.

COWLING, T. G.

1953. In Kuiper, ed., The sun, p. 532. University

of Chicago Press.

ELSASSER, W. M.

1955. Amer. Journ. Phys,, vol. 23, p. 590.

ROMPB, R., AND STEENBECK, M.

1939. Erg. Exakt. Natur., vol. 18, p. 257.
Hydromagnetic and Plasma Problems1

By Eugene N. Parker

An ionized gas or plasma possesses two

obvious general classes of dynamical modes as

a consequence of its being a mixture of two

electrically coupled gases. In the electron gas

the velocity of sound is much higher than in

the other constituent, the ion gas, so that if

the plasma is deformed sufficiently rapidly, the

ion gas will lag behind the electron gas; this

will result in separation of charge and in large

electrostatic forces that tend to restore the

initial relative configuration. The resulting

oscillations of this "electric" mode are called

plasma oscillations. In the absence of magnetic

fields, the oscillations have a fundamental

angular frequency

&>,=(4wieTi/m)>*, (1)

where e is the charge of an electron in electrostatic units, m is the electron mass in grams,

and n is the number of electrons per cm8. The

frequency defined in equation (1) is called the

plasma frequency, and we might add that in

the presence of magnetic fields there are three

possible frequencies that are algebraic combinations of wp with the electron gyrofrequency

Be/me (Westfold, 1949).

If the disturbances initiated in a plasma are

slow compared to the plasma frequency, then

obviously the electron and ion gases will move

together. In this case the plasma is well

approximated by a conducting classical fluid

(Elsasser, 1954, 1955a). Since the motions

of an electrically conducting fluid interact

with magnetic fields, and in astronomical

phenomena are so generally associated with

magnetic fields, this plasma mode may be

termed the "magnetic" mode. The study of

the dynamics of conducting fluids in the

1

 Assisted In part by the Office of Scientific Research and the OeoPhystcs Research Directorate, Air Force Cambridge Research Center,

Air Research and Development Command, V. 8. Air Force.

'Enrico Fermi Institute for Nuclear Studies, The University of

Chicago, Chicago, m.

presence of magnetic fields is variously referred

to as hydromagnetics, magnetohydrodynamics,

and fluid electrodynamics.

Hydromagnetics

The importance of hydromagnetics in astrophysical phenomena may be seen from the

fact that much of the region within the galaxy,

including stellar atmospheres and interiors, is

occupied by gas in various degrees of ionization

and readily transports electric charge. Both

direct observation and theoretical inference

now generally support the belief that much of

interstellar space, together with a large fraction

of the stars, is permeated by magnetic fields

with an energy density comparable to the

kinetic energy density of many of the gas

motions. Thus, the magnetic field stresses

are comparable to the Reynolds stresses. The

high conductivity of the matter indicates that

it is strongly coupled to the pervading field,

so that in treating most fluid motions of astrophysical interest, instead of the hydrodynamic

equation,

(2)

one must use the hydromagnetic equations,

dv

df (VXB)XB. (4)

Here c represents the conductivity of the fluid

and M the permeability in mks units.

The complexity of the nonlinear hydromagnetic equations indicates that the general

problem must be divided into the several basic

physical processes; consequently, the study of

hydromagnetics involves the investigation of

individual effects, one at a time. Instead of a
60 SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS VOL.1

formal general solution of equations (3) and

(4), we will ultimately achieve a general understanding based on quantitative understanding

of the many basic concepts (Elsasser, 1955b).

To mention but a few of the physical principles, let us consider first a fluid of very high

conductivity. Then equation (3) reduces to

f=VX(vXB). (5)

From equation (5) it is readily shown that the

lines of force move exactly with the fluid.

Formally, if C represents a contour moving with

the fluid, and S the area enclosed by the contour,

then

(6)

The lines of force retain their identity, since any

two elements of fluid threaded by the same line

of force will remain so connected. Equation

(5) is the same as the Helmholtz vorticity

equation for an inviscid fluid. Cauchy's integration gives B as an algebraic function of the

derivatives of the Lagrangian coordinates of the

fluid particles (Brand, 1947).

If the conductivity is not essentially infinite,

then of course the lines of force slip relative to

the fluid. Besides the slipping, however, they

lose their identity by continually changing

their pattern of connection. Thus, for instance,

a loop in a long tube of flux may detach itself

from the tube, giving a magnetic ring and the

original, more or less straight tube. This

detaching process forms one of the basic steps

in the regenerative hydromagnetic dynamo

discussed below.

The forces exerted on the fluid by the magnetic field are derivable from the magnetic

stress tensor, BtB}/n, (i, j=l, 2, 3), representing

tension along the lines of force, and the isotropic

magnetic pressure, B2/2fi (Cowling, 1953), in

mks units. Thus the total pressure isp-\-B2/2pi.

Within a region of magnetic field the gas pressure p will tend to be less than it is outside by

the amount B2/2n. The gas density tends to

be less within the field, resulting in the phenomenon of magnetic buoyancy (Parker, 1955a).

The tension in the lines of force and the mass

of the fluid clinging to them indicate that the

lines are analogous to stretched strings, allowing

the propagation of transverse waves with a

velocity whose square is equal to the stress

density B2\\i divided by the material density

(Alfv6n, 1950). These hydromagnetic waves

are transverse, and for this reason are generated

more copiously by the transverse fluid motions

of subsonic hydrodynamics than are the

longitudinal acoustical waves. As the magnetic energy density becomes large compared

to the kinetic energy density of the fluid, all

hydromagnetic motions reduce to hydromagnetic waves (Parker, 1955b). It is interesting

to note, however, that the presence of the

magnetic field may enhance the generation of

acoustical waves (Kulsrud, 1955).

The fundamental problem of the convective

stability of a rotating conducting fluid in the

presence of a field has been solved by Chandrasekhar (1953b; see also 1952, 1953a, 1954a)

with remarkable results. Simple Benard cells

do not occur. Instead, competition develops

between the rotational and magnetic effects

and results in two unstable regimes. One

might say, very roughly, that the magnetic

field tends to stiffen the fluid, producing

stability, and the Coriolis forces tend to produce instability. The results of Chandrasekhar's analytical formulation of this very difficult problem have been given experimental

confirmation (Nakagawa, 1955).

Chandrasekhar (1955) has also treated the

problem of hydromagnetic turbulence and finds

that there are two statistical modes, with either

the fluid motions or the magnetic field dominating at large wave numbers. The analysis

shows that there is not a close approach to

equipartition of energy between the fluid motion and the magnetic field, but, on the other

hand, it does indicate that K pv

2

 and B2/2n will

not differ by an order of magnitude.

The origin of the stationary dipole field of

earth and that of the migrating solar fields

resulting in sunspot formation are beginning

to be understood on the basis of hydromagnetic

dynamo effects whereby two magnetic fields in

a region are each generated from the other by

fluid motions. The so-called toroidal field is

produced by simple shearing from the poloidal

field; the latter is then regenerated by C37clonic
NEW HORIZONS TN ASTRONOMY 61

motions in the fluid, forming loops from the

toroidal field (Parker, 1955c; Elsasser, 1955b).

Presumably a sunspot is formed from the solar

toroidal field through a combination of magnetic buoyancy, hydrostatic equilibrium (Parker, 1955a), and the observed cooling of

the interior of the spot, though most of the

details remain to be treated quantitatively.

Probably nowhere does one find more phenomena whose origin so obviously cannot be

explained on hydrodynamic grounds than in

the violent phenomena of solar activity, such

as spicules, active prominences, flares and

surge prominences, the ejection of ion clouds,

and cosmic ray bursts. The origins of these

phenomena are today not at all understood,

but one presumes that at least some of them

admit of a hydromagnetic origin. The terrestrial consequences make solar activity a

particularly pressing challenge to hydromagnetic research.

Understanding the terrestrial effects is important, but it has often been suggested that

solar activity may prove also to be the "little

brother" to some of the grander phenomena

that result in the Wolf-Rayet stars, flare stars,

and Babcock's magnetic variables where prominence activity, flares, and sunspots have developed to pathological proportions.

Solar activity apparently ejects field-bearing

clouds of gas into interplanetary space. At the

earth the clouds are somehow responsible for

the aurorae, magnetic storms, certain ionospheric disturbances, the configuration of the

terrestrial weather pattern, and the observed

interplanetary gas density of as much as 500

hydrogen atoms per cm3 (Storey, 1953; Siedentopf, Behr, and Elsasser, 1953; Behr and

Siedentopf, 1953). Further, cosmic ray intensity variations not directly due to the

immediate emission of new particles by the

sun are now generally regarded (Neher, 1955)

as being due to the antics of interplanetary

clouds. It becomes important, therefore, to

investigate the dynamics of such clouds. If at

the same time the dynamics of ejection from

the sun can be expounded, we will be able to

check our ideas by both solar and terrestrial

observations.

Cosmic rays are presumed to be accelerated

by interaction with the magnetic fields carried

by moving matter. The acceleration is presumed to take place in the galaxy and solar

flares. One consequence of the high conductivity of most matter is that the electric field in

the frame of reference moving with the fluid is

very small. Thus, it is readily shown that

E=-YXB , (7)

where v is the fluid velocity, and there can be

no large electrostatic acceleration effects. The

general solution of the motion of a charged

particle in slowing varying hydromagnetic

fields, as well as calculations of several cases of

abruptly varying fields, indicate that there are

no hydromagnetic accelerating mechanisms

other than Fermi's mechanism (Fermi, 1949)

and the betatron effect (Swann, 1953; Riddiford

and Butler, 1952). Since both mechanisms

require that the particles initially have energies

of the order of 0.5 Bev per nucleon to avoid the

enormous ionization losses at low energies, one

presumes that there is some more vigorous

mechanism capable of accelerating particles

from thermal energies in spite of the tremendous

energy losses. Finding such a mechanism is

obvious^ of fundamental importance.

The assumption that cosmic rays are accelerated throughout the galaxy (Morrison, Olbert,

and Rossi, 1954; Unsdld, 1955; Parker, 1955b;

Elsasser, 1955a) presents energetic difficulties.

Davis (1955b) has suggested that the gas ejected

from the sun forms a cavity in the magnetic

field of the galaxy capable of trapping particles

in the solar system, though the mechanism

apparently cannot account for the cosmic ray

cutoff with the attendant correlation with the

sunspot cycle. A galactic origin of cosmic rays

and the sweeping away of galactic cosmic rays

by outgoing magnetic interplanetary clouds

seems to give a suitable cutoff, but so little is

known about the nature of interplanetary clouds

and Davis' cavity that the details of the mechanism are entirely speculative.

Besides the evidence from the polarization of

starlight (Davis and Greenstein, 1951), it has

been shown (Chandrasekhar and Fermi, 1953)

that the observed properties of the spiral arms

of the galaxy require a magnetic flux density of

about 10"5 gauss along each spiral arm. A

magnetic field of this strength permeating the

381014—56 «
62 SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS VOL. 1

interstellar medium will be important in determining the characteristics of stars formed by

gravitational collapse of the interstellar medium,

for, although it does not directly affect Jeans'

instability criterion (Chandrasekhar, 1954b),

the field will strongly influence the nature of the

fluid motions against which work the gravitational forces of collapse. Hydromagnetics evidently plays an important role in the formation,

structure, and evolution of galaxies, though

little has been done on the integrated problem.

The pulsations of some variable stars suggest

(Chandrasekhar and Fermi, 1953; Chandrasekhar and Limber, 1954; Ferraro, 1954) that

the general magnetic fields of the individual

stars play an important role in determining the

form and period of the oscillations of the stellar

body. The tremendous general stellar field

that is required implies a hydromagnetic dynamo

of immense power.

Plasma problem

The extraordinary number of modes of oscillation of the electron gas in a plasma is exhibited

in Bailey's (1950, 1951) general electromagnetoionic theory. Starting with a plasma in the

presence of a magnetic field, and using Maxwell's equations, the equation of continuity,

etc., neglecting damping by collisions, and considering small sinusoidal deviations from uniformity, Bailey finds a number of steady state

modes with frequencies that are simple algebraic combinations of the plasma frequency uP

and the gyrofrequency <ae, which is the angular

velocity of an electron in the given magnetic

field. Besides these usual modes, however, there

are modes which grow or decay exponentially by

energy transfer to the over-all drift velocity of

the ions. The presence of the magnetic field

enhances the exponential growth effects. Since

the theory neglects the second order terms, one

does not know to what extent the exponential

growth is limited.

The response of a plasma to electromagnetic

waves results in an index of refraction that can

be real, imaginary, or complex, and the propagation of electromagnetic waves through a nonuniform plasma is complicated and of fundamental importance (see Pawsey and Smerd,

1953). Consider the solar atmosphere in the

absence of magnetic fields. The plasma frequency decreases outward so that escaping

radio waves of low frequency can only have

originated in the outer corona. Higher frequencies may originate farther into the atmosphere. The observed cutoff frequency at the

low end of the frequency spectrum affords an

indication of the level of the source in the solar

atmosphere; the time rate of change of the cutoff frequency is an indication of the velocity of

the source. But, at the same time, the refraction of the waves in the solar atmosphere

complicates the problem of determining the

heliographic position of the source.

The presence of a magnetic field further

complicates the propagation, producing, among

other things, a tendency of the transverse

electromagnetic waves to propagate parallel to

the magnetic field, as a consequence of the

electrons' circling around the field. This effect

appears in the terrestrial phenomena of atmospheric whistlers investigated by Storey (1953).

The propagation depends upon the interplanetary gas density near the earth and so gives

information on the interplanetary medium.

The anomalous sources of radio noise in the

sun and elsewhere are often associated with high

velocity gas or particle streams. The rate of

decrease with time of the observed frequency of

a given solar burst suggests that the excitation

might be due to clouds that are moving out

through the solar atmosphere with velocities of

the order of 600 km/sec (Pawsey and Smerd,

1953); intense galactic sources have been identified with galaxies in the process of passing

through each other (Baade and Minkowski,

1954). The radio emission presumably arises

from the collisions of the interstellar gas clouds

contained in the interpenetrating galaxies

(Greenstein and Minkowski, 1954). The presence of high velocity streaming and the general

difficulty with thermal mechanisms (requiring

temperatures up to 1O150 K) in the more intense

sources suggest the investigation of vigorous

macroscopic phenomena such as the exponentially growing plasma oscillation (Bailey, 1950,

1951; Haeff, 1948, 1949; Pierce, 1948; Bohm

and Gross, 1949, 1950; Buneman, 1950; Malmfors, 1950; Blum, Denisse, and Steinberg, 1951;

Twiss, 1951). To establish whether plasma
NEW HORIZONS IN ASTRONOMY 63

oscillations are responsible for the intense

sources, one will have to obtain a detailed

knowledge of the streaming which supposedly

excites the oscillations. Then one must investigate the effects of the nonlinear terms and

the actual boundary conditions found in nature

in order to determine to what extent the oscillations will grow and whether the electromagnetic radiation generated by the radiations can

escape. The question of escape has already been

investigated but with conflicting results.

Perhaps the acceleration of cosmic rays is

closely associated with the origin of radio

noise; both might seem to require, or are

observed to have, association with violent

fluid motions. Hydromagnetic processes seem

unable to furnish the injection energies needed

for Fermi's mechanism and for the betatron

effect. Thus, obviously, investigations of violent plasma phenomena should be carried out

with cosmic rays as well as radio noise in mind.

But in contrast to this point of view is the

recent discovery that the visible and radio

spectrum of the Crab Nebula can only be explained as synchrotron radiation from electrons

with energies of the order of 1 Bev, which by

definition are cosmic ray particles (Shklovsky,

1953). The radio emission from the galactic

halo seems to require a similar mechanism.

Thus, because we are ignorant of the "basic"

energy source, cosmic rays become the apparent

prime mover.

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1950. Coemical electrodynamics. Oxford University Press.

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1964. Astrophys. Journ., vol. 119, pp. 206, 215.

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1950. Phys. Rev., vol. 78, p. 428.

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1950. Nature, vol. 165, p. 474.

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1952. Philos. Mag., vol. 43, pp. 501, 1317.

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1953. Astrophys. Journ., vol. 118, pp. 113, 116.

CHANDRASEKHAR, S., AND LIMBER, D. N.

1954. Astrophys. Journ., vol. 119, p. 10.

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1953. In Kuiper, ed., The sun, p. 532. University

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DAVIS, L., JR.

1955a. Summer meeting of the American Physical

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1955b. Phys. Rev., vol. 100, p. 1440.

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ELSASSER, W. M.

1954. Phys. Rev., vol. 95, p. 1.

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1954. Astrophys. Journ., vol. 119, p. 407.

GREENSTEIN, J. L., AND MINKOWBKI, R.

1954. Astrophys. Journ., vol. 119, p. 238.

HABTT, A. V.

1948. Phys. Rev., vol. 74, p. 1532.

1949. Phys. Rev., vol. 75, p. 1546.

KuLBRUD, R.

1955. Astrophys. Journ., vol. 121, p. 461.

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1950. Ark. Fys., vol. 1, p. 569.

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1954. Phys. Rev., vol. 94, p. 440.

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1955. Nature, vol. 175, p. 417.

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1955. Spring meeting of the American Physical

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64 SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS

RlDDIFOBD, L., AND BtJTLEB, S. T. SWANN, W. F. G.

1952. Phuos. Mag., vol. 43, p. 447. 1933. Phys. Rev., vol. 43, p. 217.

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Solar System

Jiartti s Magnetism

By Walter M. Elsasser

 1

An explanation of the phenomena of geomagnetism is of interest to theoretical astrophysics not just because the earth happens to

be a celestial body, but mainly because to

reproduce such phenomena in a conventional

laboratory is extremely difficult if not nearly

impossible. Physical processes in the universe

clearly fall into two classes: those that have a

readily accessible laboratory equivalent (e. g.,

spectroscopy, nuclear reactions), and those that

do not (e. g., radiative viscosity, large-scale

turbulence). The phenomena we choose to

designate as hydromagnetic, of which geomagnetism is the prime example, belong to the

second class. The nearest hydromagnetic phenomena beyond the confines of the earth are

those on the sun, particularly in sunspots, but

their study has shown them to be exceedingly

complex.

Babcock's now classical investigations of

stellar magnetic fields started as a search for

magnetic fields in early-type stars. These were

chosen because many of them are known to

rotate rapidly and a connection between rotation and magnetic fields was suspected. The

conditions under which stellar magnetic fields

can be ascertained are extremely stringent: the

stars must be bright enough (>7th mag.), and

the average magnetic field over the visible face

of the star must be at least several hundred

gauss. Of the stars investigated, 65 had to be

rejected as spectroscopically inadequate for the

analysis, 35 yielded measurable magnetic fields,

and in another 20 magnetic fields are suspected

but not yet assured (as of 1953, the date of

Babcock's review). As Babcock points out,

this is an extraordinarily large yield even from

a purely statistical viewpoint, and the presumption is justified that very many stars possess

magnetic fields of appreciable magnitude.

There is also reason for the belief, based on the

> Department of Physics, University of Utah, Salt Lake City, Utah.

investigations of A. Deutsch, that phenomena

equivalent to sunspots, but sometimes of much

larger extent, occur in many stars. Perhaps it

is useful to recall that the magnetic field is the

dynamically essential feature of a sunspot, as

evidenced by the fact that the mechanical

forces exerted by the magnetic field exceed all

other forces acting in a sunspot.

The main interest in cosmic magnetic fields

centers on two problems: first, the way in

which they influence observable astrophysical

phenomena, stellar atmospheres, interstellar

gas motions, radio noise, the acceleration of

cosmic rays, and others; and second, the manner

of their generation. The first class of problems

lies outside the domain of this paper, but is

dealt with elsewhere in this volume. The

earth does furnish us with a prime example of

the mechanism of generation of large-scale

magnetic fields; indeed, there is reason to believe that geomagnetism lends itself somewhat

more readily to the construction of a consistent

theory than do the phenomena in the solar

convection zone which give rise to sunspot and

solar magnetism, not to speak of the magnetic

fields of stars whose details are barely, if at all,

accessible to observation.

The ubiquitous presence of magnetic fields

in the universe leads one to accept a model of

the generation of the earth's magnetic field

which demonstrates that magnetic fields can be

generated, and maintained, by inductive amplification in an electrically conducting fluid of

sufficiently large dimensions. This is the

dynamo theory for which a great deal of observational evidence now exists. In keeping

with the nature of this volume we shall not

dwell on the evidence in detail, but shall

emphasize the basic dynamical features and

their implications.

The geomagnetic dynamo is located in the

earth's liquid core (diameter 0.55 of that of the

67
68 SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS

earth). There seems little doubt that this core,

which seismological evidence has shown to be

fluid, consists of molten iron. It may be shown

by an analysis of the observed field that the

sources of the geomagnetic secular variation are

located near the boundary of the core. The

geomagnetic secular variation is a very remarkable phenomenon; it is the only effect originating in the interior of the earth whose spectrum

has periods of order of some tens up to some

hundreds of years, whereas all geological processes are of course measured in millions, at the

best in hundreds of thousands of years.

The earth's magnetic field (if we ignore the

very small component that originates in the

ionosphere) consists of a dipole inclined by

11K° relative to the earth's axis, and of an

irregular part made up of all the higher terms

of a spherical-harmonic analysis. If the time

dependence of the irregular field is considered,

it is found that in all probability this irregular

field does not have a component constant in

time. The dipole axis also changes in longitude

and latitude, but much more slowly than the

higher spherical harmonics. There is at present

a good deal of evidence from paleomagnetism

(the study of the remnant magnetization of

rocks) to the effect that the present angle

between the geographical axis and the magnetic

dipole axis is larger than it was in the (r. m. s.)

average over past geological periods, and that

the statistical distribution of these angles is

centered about the geographical axis. This

result indicates the reality of a close connection

between the earth's rotation and its magnetic

field which the dynamo theory postulates.

The secular variation (which is essentially

the time derivative of the nondipole component

of the geomagnetic field) originates in a relatively thin layer at the top of the core. The

reason for this is that a layer of metal acts as a

screen for time-dependent fields, and the periods

of the secular variation are such that they would

be effectively screened off by a layer, say, 100

km thick. Thus the secular variation cannot

give us information about the interior of the

core where the essential part of the dynamo

mechanism is located, but it does teach us

something about the outer layers of the core

and, incidentally, about the relation of the core

and the solid part of the earth, the mantle,

surrounding it.

In order to explain the observations it is

necessary to assume that the mantle has a small

electrical conductivity. This is plausible considering the temperatures of perhaps 2,000-

4,000° prevailing in the lower mantle. Some of

the magnetic field generated in the core on

penetrating into the mantle sets up a mechanical force there. As a result of the action of

this force, there is a small difference in angular

velocity between mantle and core, the mantle

rotating slightly faster than the top layers of

the core. In the stationary state a balance is

achieved between the accelerating force just

mentioned and an opposing eddy friction, the

latter probably also of a magnetic nature. The

differential rotation has been observed and

analyzed; it is known as the westerly drift

because for an observer stationed at the earth's

surface, the individual features of the irregular

magnetic field seem to undergo a slow average

motion toward the west. It has also been

shown that this motion is suffering fluctuations

of moderate size over periods of 10-20 years.

These fluctuations appear correlated with the

fluctuations in the earth's (more precisely, the

mantle's) rate of rotation which are directly

observed in positional astronomy of the moon

and planets. Theoretical analysis has led to

the result that these fluctuations in the earth's

angular velocity can be satisfactorily explained

in terms of a slightly variable magnetic coupling

between core and mantle. The required fluctuations of the magnetic fields near the boundary of the core are quite in keeping with what

we might expect on the basis of our general

knowledge of the earth's magnetic field. The

fluctuations may be either in the over-all field

producing a torque on the mantle or may be

due to changes in the (magnetic) eddy friction

which counterbalances the accelerating torque.

In any event, this explanation of an important

astronomical phenomenon seems now to rest on

a sound geophysical basis.

As we go into the interior of the core there

appears a type of magnetic field which has no

counterpart on the outside, the so-called toroidal

field. In a pure toroidal field the field lines

would be circles about the earth's axis (the
NEW HORIZONS IN ASTRONOMY 69

actual field in the earth is a superposition of

more than one type and hence more complicated). The toroidal field, since it cannot be

observed directly, had to be inferred from solutions of the electromagnetic field equations,

but its existence does not underlie any serious

doubt since its existence is mathematically necessary. The toroidal field apparently is appreciably larger than the dipole field in the interior

of the core, being perhaps of order 20-40 gauss.

Its counterpart also exists in the solar convection zone, and there is reason to believe that it

provides the field from which the sunspot fields

are ultimately formed.

Some 20 years ago C. T. Cowling showed that

a fluid in which the motion is confined to meridional planes cannot form a dynamo. This result may be generalized to the statement that

no fluid motion of rotational symmetry about

an axis can maintain a dynamo. There are good

theoretical reasons, furthermore, to believe that

no two-dimensional motion, that is, no motion

in which the fluid particles are confined to surfaces, is capable of providing dynamo action.

Hence dynamos have an essentially three-dimensional geometry; this makes the structure

and dynamics of even the simplest models so

complex that we prefer not to expound them

here in any detail. Instead, we shall confine

ourselves to a discussion of the essential ingredients of dynamos, the general theorems on

which their operation is based and which make

their existence plausible. So far as we can judge,

these principles are most likely to be the same

for all hydromagnetic dynamos, the terrestrial

one as well as the solar and stellar ones.

The simplest formal conditions appear in the

limit of infinite electrical conductivity. One can

then show that

ItJ B n d<r= (1)

a theorem which states that the magnetic flux

normal to a given arbitrary surface remains constant in time as this surface is carried with the

fluid. This result, first given by Cowling, has

sometimes been expressed by saying that the

magnetic lines of force are "frozen" in the fluid

and are bodily carried along by it. It may

readily be seen, then, that the magnetic field can

be sheared or twisted locally by the fluid motion

so that amplification results. It remains to be

demonstrated that amplificatory processes can

be organized in such a way that they can maintain the magnetic fields observed on the earth,

sun, and stars. This turns out to be an exceedingly difficult mathematical problem. Those

who seek a formal solution after the manner of

existence proofs for the solutions of differential

equations are likely to remain disappointed.

The problem is essentially nonlinear (clearly, no

linear system can exhibit self-sustained amplification) and our understanding of the behavior

of nonlinear dynamical systems is still in an extremely rudimentary state. Those making an

analysis must, for the time being, be content

with inadequate rigor and must rely to some

extent on intuition, keeping in mind that there

are definite phenomena waiting for explanation.

The present writer finds this situation rather

challenging to the theoretician, but he realizes

that it does not always appeal to the more

mathematically inclined analyst.

Equation (1) applies only to a medium of

infinite conductivity. In such a fluid we could

not have a steadily operating dynamo because

the continued stretching and twisting of lines

of force would (as Bondi and Gold have

remarked) result only in a progressive confusion

of magnetic fields, not in a stable pattern.

In a fluid of finite electrical conductivity, the

right-hand side of (1) is no longer zero. The

ratio of the right-hand to the left-hand side

is then measured by a nondimensional quantity

known as the "magnetic Reynolds number,"

T» T TT (n\

where L and V are length and velocity, used in

the manner known from the conventional

hydrodynamic Reynolds number, and 0- is

the conductivity of the fluid (in e. m. u.).

Rm is large, that is, the conservation theorem

(1) is approximately fulfilled when we are

dealing with large linear dimensions. In the

core of the earth and still more in the stare,

Rm is as a rule quite large numerically. This

fact can be expressed in a rather concrete

physical language by saying that for large

Rm, the time required for spontaneous decay

of the electromagnetic fields in the fluid is
70 SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS

long compared to the periods of the mechanical

fluid motion; Rm as defined by (2) can be shown

to be indeed the ratio of these two time constants. In the laboratory Rm is as a rule numerically small; this cannot readily be remedied

since all three factors on the right of (2) are

clearly limited in laboratory experiments.

Since the functioning of hydromagnetic dynamos depends essentially on the field's being

carried along by the fluid, that is on large

Rm, it follows that hydromagnetic dynamos

are essentially a matter of cosmic physics;

namely, of very large linear dimensions. This

substantiates our previous contention that

the earth is the nearest thing to a laboratory

of bydromagnetism, at least so far as dynamos

are concerned.

If the magnetic lines of force are continuously

stretched and twisted, as is the case in the

turbulent regime of a conducting fluid, the

average result may be shown to be amplification of the existing field. This increase is

likely to stop when a state of equipartition is

achieved where there is an amount of average

magnetic energy equal to the average kinetic

energy of the fluid motion. Hydromagnetic

turbulence fields are of great astrophysical

interest (they have been investigated by

Chandrasekhar) but it is clear that they do not

constitute dynamos. The latter are called on

to produce and maintain observed magnetic

fields which are of a high degree of regularity;

e. g., the earth's dipole field.

We have mentioned before that a fluid can

act as a hydromagnetic dynamo only if its

motions are essentially three-dimensional. Thus

a stably stratified fluid where motions can at

best take place in the approximate direction of

equipotential surfaces cannot act as a dynamo.

Under the actual conditions of geophysics and

astrophysics the fluid must be thermally

unstable and hence in convective motion.

We do not know the exact causes of convection

in the earth's core. We are better off with

regard to the sun and stars; there we can be

rather certain that the hydrogen convection

zone is the seat of amplifying processes that

lead to hydromagnetic dynamos.

A relation between rotation and dynamo

action is clearly indicated by the behavior of

the earth's magnetic field and is corroborated

by much astrophysical evidence. Theoretical

arguments lead to the conclusion that rotation

of the conducting fluid is essential for dynamo

action. If a convective layer rotates, the convective transport of angular momentum leads

to an average state where the layer rotates

nonuniformly, the outer parts of the layer

rotating less rapidly than the inner. If we

assume a dipole field in the interior of the

earth's core, such a field will be drawn out by

the nonuniform rotation so as to give rise to a

toroidal field. The toroidal field in turn sheds

magnetic eddies which, under the influence of

the Coriolis force, may be shown to behave in

such a way that reinforcement of the "poloidal"

(dipole) field results. Through the theoretical

investigations of recent years it has become

very probable that dynamo action results from

the interplay of mutual amplification of two

types of field components, the toroidal and

poloidal fields. They correspond roughly to

the two coils, say the stator and rotor of a

conventional electrical generator. In the analysis of the geomagnetic dynamo this mechanism

appears fundamental; the stellar dynamos are

perhaps somewhat more complicated, but

ultimately should reduce to a similar type of

interaction.

A conspicuous feature of dynamo mechanisms of this type is their low geometrical

symmetry. Convection excludes any motions

that do not have a radial component, and the

Coriolis force which is not invariant with

respect to the interchange of east and west

leads to patterns of fluid motion that do not

have rotational symmetry about the earth's

axis. The fluid motions required for the terrestrial dynamo do not seem to have any

geometrical symmetry, and this is probably

characteristic of hydromagnetic dynamos in

general. Even in technical generators the

system of wires, commutators, etc., forms a

low-symmetry structure. In this requirement

of low symmetry, hydromagnetic dynamos

differ from most of the traditional problems

of classical dynamics where it has been the

custom to make problems accessible to solution

by using symmetry operations for their mathematical simplification. This situation might

account for the slow pace with which the

theoretical explanation of the phenomena of
NO.I NEW H0RIZ0N8 IN ASTRONOMY 7 1

terrestrial and solar magnetism has developed References

in the past, in spite of the fact that many of the

observations are far from recent. It is hardly We have omitted a list of detailed references

necessary to say that under these conditions since they may be found in the following fairly

the subject offers many challenging opportu- extensive reviews:

nities for analytical investigations of funda- ELSAS8EB ( WALTB B M .

mental importance to astrophysics. The going 1950 Re v Mo d Fhy8> vo l ^ p L

will admittedly be difficult and new analytical ^55 Amer. Journ. Phys., vol. 23, p. 590.

techniques seem to be called for before one 1956a. Rev. Mod. Phys., vol. 28, p. 135.

can hope for major successes. 1956b. Amer. Journ. Phys., vol. 24, p. 85.

Time

By William Markowitz1

The problems of chief interest concerning

time are the precise nature of the variations in

the speed of rotation of the earth and the relation between the gravitational and atomic

time scales.

The first problem deals with the basic properties of the earth. The second deals with a fundamental physical property of the universe. Although we do know that variations in the speed

of rotation of the earth occur, we are not certain

of the details. In fact, within the last few years

many of the ideas on the subject have been drastically revised. As to the second problem, we

have no idea at all, at present, whether or not

the gravitational and atomic time scales are

the same; this must be decided by experiment.

It has been known for some time that the

speed of rotation of the earth is subject to

changes, which may be classified as periodic, irregular, and secular. These variations made

the mean solar second unsatisfactory as the

fundamental unit of time for scientific and technical purposes. Accordingly, in September 1955

at Dublin, the International Astronomical Union redefined the unit of time as a specified

fraction of the tropical year for 1900.0. The

new unit is identical with the second of Ephemeris Time, and is obtained in practice from observations of the moon.

I shall briefly describe recent developments

concerning studies of the rotation of the earth

and the instruments and clocks used in determining time.

The yearly and semiyearly periodic terms

were found by Stoyko at Paris in 1937 and were

confirmed at Greenwich. Later, large changes

in these terms from year to year were indicated.

However, a homogeneous solution based on observations made with the photographic zenith

tubes (PZT's) of the U. S. Naval Observatory

(Markowitz, 1955) showed little change from

1

 Director Time Service, T7. S. Naval Observatory, Washington, D. C.

year to year for the interval 1951-1954. Shortperiod terms due to tidal action of the moon

were also found.

It was formerly believed that sudden changes

in speed of rotation as large as several milliseconds per day occurred at irregular intervals.

Brouwer (1952), however, has been led to the

hypothesis that the irregular changes are caused

by small, cumulative, random changes in speed

of rotation and that the large sudden changes

do not occur.

Tidal friction tends to slowly diminish the

speed of rotation. Holmberg (1952) has recently pointed out that atmospheric tides may increase the speed of rotation, an idea due to

Lord Kelvin. Hence, we do not know whether

the earth is slowing down or speeding up.

The quartz-crystal clock, which has replaced

the pendulum clock in precise timekeeping, had

been developed by 1950 to the point where it

could reliably be used to determine the periodic

terms in the speed of rotation of the earth. In

1955 the announcement was made that an

atomic standard of high precision had been constructed which can be compared with astronomical observation (Essen and Parry, 1955).

A program is now under way to obtain the

frequency of the atomic standard in terms of the

Ephemeris second as determined at the U. S.

Naval Observatory with the dual-rate moon position camera. If the frequency changes in the

course of time, the gravitational and atomic time

scales are not alike. An answer may be obtained within 10 years.

The atomic standard will also be used to

study variations in the speed of rotation of the

earth by comparing Universal Time with atomic

time.

Thus, we shall shortly be comparing atomic

time with astronomical time—Ephemeris and

Universal—to provide information on the problems cited. The atomic standard apparently

73
74 SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS VOL.1

reproduces the behavior of the cesium atom to

one part in 109, and higher accuracies will

probably be attained. At present the frequencies that can be obtained from astronomical

observation have precisions of several parts in

109. The probable error of the differences

between atomic time and astronomical times

will be due ahnost entirely to uncertainties in

the latter. Hence, it is the astronomical determination of time which must be improved. An

artificial satellite which does not move within

the earth's atmosphere could provide another

determination of Ephemeris Time. We must

wait and see, however, what particular satellites

are created and how well the position can be

determined. In general, we must look forward

to continuing and improving the present techniques in time determination.

We shall now consider what needs to be done.

The Committee on Needs in Astronomy asks

where financial aid might prove most beneficial

in aiding progress. Dr. Struve (1955), in his

article "The General Needs in Astronomy,"

asks a more basic question; namely, should

more attention be given to the type of positional

astronomy carried out at the U. S. Naval

Observatory, and should more research be done

in this country with respect to the earth as a

body?

I should like to answer these questions for

the fields of meridian astronomy and for time

determination. Both types of observation provide basic astronomical data which can not be

obtained in any other way. Hence, research

in both fields should be encouraged in order

that results of the highest possible precision

may be obtained.

The determinations of time and of fundamental positions have become highly specialized

fields, and are carried on in the United States

only at the Naval Observatory. The observing

techniques have been improved over the years,

and other astronomers in the United States

have been relieved of the necessity of doing

similar work. On the other hand, these astronomers have become isolated from developments in these fields. In the United States

instruction is not given at any school or observatory in time determination or meridian

astronomy as currently practiced. The end

result is that, because of the general shortage

of scientists and of astronomers as a whole,

there is a particularly acute shortage of astronomers entering these fields.

For the past few years the Naval Observatory

has not been able to fill vacancies in the Time

Service and Transit Circle Divisions with astronomers, and the outlook for the future is not

encouraging. What is desired is that enough

astronomers be trained so that work in these

fields can be actively pursued at the Naval

Observatory, and so that, in addition, research—

at least in theoretical aspects of the problems—

can be pursued at some other observatories or

schools.

From time to time foreign astronomers and

students have come to the Naval Observatory

for study and research. Such visits could also

be made, with considerable benefit, by American

astronomers and students. The minimum useful period would be about 2 weeks, while for

research and thesis purposes at least a semester

would be required.

Hence, if money is to be spent to promote

advances in these fields, it can best take the

form of grants for study at the Naval Observatory. Grants could be made to teachers and

writers of textbooks in astronomy and even to

a few undergraduates. At the present time

money is available for research grants, but few

requests for research in these special fields will

be received until there are more astronomers

able to use the grants.

References

BBOUWBH, D.

1952. Astron. Journ., vol. 57, p. 125.

ESSEN, L., AND PARRT, J. V. L.

1955. Nature, vol. 176, p. 280.

HOLMBEKG, E. R. R.

1952. Monthly Notices, Geophys. Suppl., vol. 6,

No. 6.

MABKOWITZ, W.

1955. Astron. Journ., vol. 60, p. 171.

STRUVE, O.

1955. Publ. Astron. Soc. Pacific, vol. 67, p. 214.
Geographical Positions

By John A. O'Keefe *

The determination of position on the surface

of the earth is the subject of two branches of

astronomy; namely, field astronomy and geodesy. The latter is not always considered as

a branch of astronomy, but Simon Newcomb

in the Encyclopedia Britannica has given us a

precedent for doing so; and it seems very

reasonable on historical and logical grounds.

It is the function of field astronomy to provide the initial orientation and position for the

surveys which determine geographical position.

Geodesy then extrapolates these positions by

means of measurements on the ground to cover

a whole country, or even a whole continent,

without further reference to the stars, except

for orientation. The role of extrapolation in

geodesy is unique, so far as I know, in science.

If it were physically possible, the geodesists

would certainly prefer to have only a single

astronomical determination for the whole world

and to rely on geodesy for all other determinations. The reason for this is, of course, that

because of the irregularity of the earth's crust,

the direction of the vertical at any point is not

exactly what would have been expected on the

basis of a perfectly ellipsoidal earth. So long

as we are thinking only of navigation at sea

these errors, which very rarely amount to as

much as one nautical mile, are insignificant.

But when we deal with precise land surveys,

we can extrapolate relative positions halfway

round the world from the measurements of

triangulation more precisely than they can be

found from on-the-spot measurements by field

astronomy.

As a consequence, vast nets of triangulation

have developed which must, for the best results,

be adjusted simultaneously. The feat of adjusting systems of equations that involve several

thousand unknowns has been accomplished

repeatedly in the last few years.

1

 U. 8. Army Map Service, Washington, D. O.

Because he cannot measure vertical angles

with anything like the precision with which

horizontal angles are measured, the geodesist

lives in a mathematical world which is exactly

like the Flatland of Abbot's famous book, so

long as it is a question of geographical positions.

He knows the lengths of his lines and the

angles between them, but he does not know the

space configuration of the surface upon which

he works—the sea-level surface and its prolongation un4er the land. The maneuvers by

which he succeeds in avoiding the unanswerable

questions about the space situation bear a

remarkable resemblance to the squirming of

the relativists when asked how their curved

space would look from the outside. The resemblance is a family one, since C. F. Gauss is

the father of both systems of thought.

Of course there are techniques for finding

heights, but these are only with relation to

this sea-level surface of unknown form.

The principal application of the techniques

of field astronomy in geodesy is therefore in

determining the form of the sea-level surface of

the earth, for each measurement of an astronomical position determines the slope of the sealevel surface at a point. The difference between the astronomically measured position and

the geographical position determined geodetically is, in fact, the slope of the sea-level surface

at that point (if we assume that the slope at

the initial point was zero). From a large number of such determinations of slope, the form

of the sea-level surface can be found by

integration.

Finally, we may also seek to find the form

of the sea-level surface of the earth from the

measurements of gravity. There is a famous

theorem, called Stokes' theorem, which connects the form of the sea-level surface of the

earth with the measured values of the intensity

of gravity.

76
76 SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS VOL.1

Thus the classical approach was to deal in

terms of known angles and known distances on

a largely unknown surface, the determination

of whose form was left to a small group of

geophysicists, since the mathematicians had

found ways to bypass the requirement for it.

By extrapolation, vaster and vaster systems

of triangulation have been constructed, until

now we have most of the Western Hemisphere

on one system, and most of the Eastern Hemisphere on another.

The new techniques

The geodimeter is rapidly replacing the invar

tape for the measurement of baselines. The

geodimeter is essentially the result of reversing

the Fizeau method of getting the velocity of

light so that, from the known value of the

velocity, the distance is found. The toothed

wheel is replaced by a Kerr cell, modulated

by ultra high-frequency voltages. The new

technique has proved remarkably successful

in trials over the last few years, since it measures

directly the triangulation sides. The old invar

tape method could be employed only on relatively short lines, which then had to be related

to the long lines of the regular chains by an

elaborate base-expansion net.

The employment of radar for the measurement of great distances is rapidly expanding.

The most useful tool at the present time is the

modification of the wartime Shoran, which is

known as Hiran. The principal new feature

of Hiran is the attempt to reduce the signal

intensity correction by the use of gain control.

The Mediterranean has been bridged by this

method from Crete to Libya; the North Sea

from Scotland to Norway; and the Straits of

Florida from Key West to Cuba.

For yet longer lines, radar techniques do

not seem to work. The methods that go over

the horizon depend on the ground wave; and

the velocity of the ground wave depends, it

seems, on the nature of the ground. Repeatability is rather easily secured; but the interpretation of the results offers considerable

difficulty at the present.

The moon is now being employed for measures over the horizon. The parallax of the

moon, unlike that of any other celestial object,

is large enough to furnish a method for the

determination of place on the earth. There

are several ways of using this fact; each of

them deals in its own way with the twin

difficulties of the remoteness of the moon and

the irregularity of its surface.

The most perspicuous method is that of

Markowitz, who simply photographs the moon

in a star-field and measures the distances from

points of the moon's limb to the stars. Simple

as the theory is, its practice was not mastered

until recently. The light of the moon must be

reduced by a filter so that it can be observed on

the same photograph as an eighth magnitude

star, 100 million times fainter. The motion of

the moon during the time required to obtain an

image of an eighth magnitude star amounts to

several miles; this must be allowed for by a

special optical device which, in effect, arrests

the motion of the moon relative to the stars for

a few seconds. The measurements are made

with great care with a precise measuring engine,

and the effects of limb irregularities are taken

out by averaging about 30-odd points. The

result of a pair of plates has a probable error of

0.15 second of arc in each coordinate, which

amounts to about 900 feet on the ground because of the moon's remoteness. To obtain

more precise results on relative positions, it is

now planned to obtain a long series of observations.

A second method, which aims to obtain

greater precision in the individual observations,

makes use of eclipses. Near the moments of

second and third contact, the relative positions

of the sun and moon may be determined from

photometric measurements of the flash spectrum or of the Baily beads. Since the limbs of

the sun and moon are alined in space, there is

no disturbance from atmospheric refraction of

this relation. In addition, the resolution

theoretically obtainable far exceeds the resolution of the best telescopes.

The weakness of this method is that the

photometric observations are disturbed by

scintillations of the beads. This is a serious

matter because the diminution of brightness at

the sun's limb is not instantaneous; hence it is

necessary to fix an arbitrary level of brightness

as corresponding to the limb. The precise

determination of the moment when this brightness is reached is troubled by the scintillation.
NEW HORIZONS IN ASTRONOMY 77

The precision of a tenth of a second in arc which

is required for the success of these methods is

thus difficult to reach even for an individual

bead. For the elimination of the irregularities

of the lunar limb, reliance is placed on the

averaging of large numbers of points.

A very similar technique starts from observations of the occultation of stars by the moon.

For precise results it is not sufficient to rely on

visual observations because of the relatively

long and uncertain reaction time of the human

nervous system. Photographic observations

are likewise excluded because of the relatively

low efficiency of the photographic plate in

detecting faint stars in a short time. The only

practical method has turned out to be photoelectric.

As compared with eclipses, occultations are

relatively sharp and well-defined phenomena,

consequently there is relatively little disturbance from scintillation. The precision with

which the angles can be measured is of the

order of 0.005 second of arc, which partly compensates for the moon's remoteness. On the

other hand, there is no chance to allow for the

lunar irregularities by averaging, since a single

occultation refers to a single point of the moon's

limb. To remove the effects of the irregularities, it has been found necessary to place

the observers in such a way that all observers

of a given occultation see the star disappear

behind the same point of the moon's limb.

This means that the telescopes must be portable and that special computations must be

made to determine the points at which they

are to be placed. Not more than two occultations can be observed per month by a single

astronomical party because of the problems of

movement and survey.

The lunar methods of determining position

do not fit into the conventional picture of the

old triangulation. They do not fix either the

direction or the distance between the points

involved; rather, they yield condition equations involving both the horizontal and the

vertical positions. They represent the approach of a new point of view in geodesy.

The geodesy of the future

A development that will most certainly come

with the future is the three-dimensional attack

on geodetic problems. In aerial photography

as a method of mapmaking, we have already

seen the entry of the three-dimensional approach into the actual preparation of the

topographic map. With the development of

aerial triangulation, we are witnessing in our

times the gradual but steady superseding of the

lower order triangulation. At present, in advanced mapping projects, only the first-order

triangulation is felt to be required, unless

there is a question of cadastral mapping.

Even for cadastral projects, there are reports

that the necessary accuracy has been attained

by aerial methods.

The most important change in the geodesy

of the future will come, perhaps, from the

application of the satellite. Its use will present a genuinely three-dimensional problem, so

that the determination of height with respect

to the center of the earth will be on a par with

the determination of ground position.

The satellite can be employed like the aircraft of the Hiran and flare observations.

That is to say, its position can be fixed from

one set of ground stations while another set of

positions in a new area is fixed from the satellite. This procedure parallels the procedure

in flare triangulation, where the aircraft position is fixed by one group of stations by intersection while the positions of the unknown

stations are fixed by resection (three-point fix)

from the aircraft. However, there is one

important difference. In speaking of positions

in the flare problem, we mean only the horizontal coordinates; in fact, the vertical coordinate is not normally measured. In the satellite case, the vertical angle will be measured

along with the horizontal angle, since the altitude above the horizon will be sufficient to

permit this; and hence we shall have what the

photogrammetrists call the problems of space

intersection and space resection.

There is another difference, however, which

arises from the fact that the satellite position

can be successfully extrapolated forward, or

interpolated, so that the observations of intersection and resection do not have to be simultaneous, as they must be in the case of aircraft.

This greatly increases the flexibility and range

of the method.
7 8 SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS VOL. a

In addition, the satellite can be used to deter- expressed in latitudes and longitudes. Almine more precisely the parameters that fix the though this question could be avoided by the

size and shape of the earth. This subject, which use of rectangular coordinates, as a practical

belongs under another heading in this book, is matter it is not; and hence we can expect that

yet germane to the discussion of geographical the establishment of definitive values of the

positions because the assumed numerical values earth's semimajor axis and of flattening will

of the size and shape of the earth are formally materially assist in the international coordinainvolved in all surveys or maps which are tion of survey and map material.
Celestial Mechanics

By G. M. Clemence1

The unsolved problems in celestial mechanics

are very numerous; in fact, relatively few problems have been solved with all the rigor, generality, and elegance that may be desired.

The problems are of widely varying degrees

of generality. In the most general case we may

study the character of the motions of any system of bodies, real or imagined, under any initial

conditions whatever. In the least general case,

we must be able to calculate the actual orbit of

a particular body for .a limited time and to

compare it with observed positions of the body;

our purpose in this case may be either to find

the body on a future occasion or to study the

discrepancies between theory and observation.

It probably is not possible to say that one

such problem is more important than another,

in an absolute sense. I have selected for discussion here two that seem especially interesting

because of the many practical applications that

could immediately be made, once the solutions

become known, and also because fast calculating

machines will be of no assistance in solving

them. Thus I may disprove the common belief

that such machines can supply all the answers to

problems in celestial mechanics.

The first may be called the problem of the

generalized planetary theory. The solar portion

of the lunar theory has been brought to a high

degree of formal perfection; the coordinates of

the moon relative to the earth, as affected by

the disturbing action of the sun, are represented

as the sum of trigonometric series, each term of

which is strictly periodic with a constant coefficient. Thus the series may be evaluated for

any value of the time whatever. In the development of the series, advantage is taken of the

smallness of the ratio of the distances of the

moon from the earth and sun, with expansions

being made in powers of this ratio.

1

 Director, Nautical Almanac, TJ. 8. Naval Observatory, Washington,

D.C.

The heliocentric motion of a planet under the

disturbing action of another planet is a very

different case, since the corresponding ratio is

so large in general that expansions in its powers

are not practicable. In the methods so far used,

the difficulty has been overcome by allowing the

time and its positive integral powers, reckoned

from an arbitrarily chosen initial epoch, to appear as factors of the trigonometric terms, outside of the arguments. Thus the terms lose their

strictly periodic character. As the time increases without limit so do the coefficients, and

the series necessarily loses its physical significance at times very remote from the initial

epoch. In the applications of the method so far

made, the practical value is limited to a few

thousand years at most.

The ability to construct a planetary theory

having the same generality as the solar' portion

of the lunar theory has long been recognized as

very desirable, and several unsuccessful attempts have been made to devise a practicable

method. The failures should not deter anyone

from taking up the problem again. The qualifications needed are a sound basic knowledge of

the principles of celestial mechanics and of the

art of practical computation, acquaintance with

the power and limitations of electronic calculating machines, and a readiness to discard the

traditional methods for new and unconventional

ones.

The second problem to be mentioned is that

of adapting the conventional planetary theory

to automatic electronic calculating machines.

By conventional planetary theory is meant the

kind already mentioned, where the time and its

powers appear as factors of the periodic terms.

In spite of its limitations the theory is of great

value and deserves many more applications

than it has had. The principal difficulty in

adapting it to automatic calculation is that

the most discriminating judgment is required

79
80 SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS VOUl

at all stages of the work if the amount of

arithmetic is to be kept within reasonable

bounds. For example, the actual calculations

involved in Hill's theory of Jupiter and Saturn,

which occupied him for 8 years, could be performed in a very few hours with a modern

machine if it were practicable to incorporate

all of Hill's judgments in the automatic program. But to devise such an elaborate program might well require 8 years of detailed

planning, with even more liability to error than

was the case with Hill's own work. The

obvious alternative, to simplify the program

at the cost of letting the machine do unnecessary calculations, is not attractive; it appears

that with the most simple program conceivable

the machine would do about a million times as

much arithmetic as is necessary, and that with

the sort of program that could be laid out in a

year the machine would still work for some

thousands of hours. Another possibility is to

start and stop the machine a thousand times

in the course of the work in order to give the

astronomer an opportunity of exercising his

judgment. But the most attractive method

would be to recast the form of the developments

completely so as to avoid the numerous multiplications of pairs of trigonometric series,

which have to be integrated term by term, and

to substitute for them a more easily controlled

numerical representation of the disturbing

forces. It may reasonably be expected, although it is not certain, that the new method

would involve more numerical work than the

old; we would willingly accept any increase up

to, say, a hundredfold.
Meteorites

By John S. Rinehart1

Meteorites are bits and pieces of extraterrestrial matter which have survived the rigors

of their passage through the earth's atmosphere.

Small ones and some quite large ones are

recovered. They are our only ponderable

extraterrestrial material. As such, they are a

storehouse of information to the astronomer.

The science of meteoritics embraces a wide

diversity of technologies: geochemistry, geology,

physics, chemistry, aerodynamics, astronomy,

astrophysics, mineralogy, and metallurgy.

The meteoriticist carries out a variety of

activities. He may simply collect, catalog,

and display meteorites. He may study their

metallurgical and crystallographic structures

and cogitate on the cosmological significance

of his findings; he may measure the relative

abundances of chemical elements in meteorites

and try to relate them to the universe as a

whole; he may investigate the craters that

have been formed by meteorites striking the

earth, or the debris that a meteorite spewed

about as it hurtled through the air.

In contrast to many other branches of

astronomy, meteoritics, in a broad sense, has

been seriously neglected by scientific investigators. Although much has been written about

meteoritics in recent years, we still lack a well

integrated concept of the science. The goal

of future studies must be to effect this integration. Before this can be accomplished, a

multitude of specific problems will have to be

defined and solved.

Where do they come from, these metallic

and stony calling cards from outer space?

What are they like? How many are there?

The answers to these questions will tell us a

great deal about the origin of the solar system,

the structure of the major and minor planets, and

the relative cosmic abundances of the elements.

' Astrophysics! Observatory. Smithsonian Institution, Cambridge,

M M .

The classification and description of meteorites is in good order. Meteorites are named

according to the locality in which they are

found. They may be classed roughly as irons,

stony-irons, and stones. The irons contain

three main constituents in varying proportions:

kamacite, which corresponds to the a-phase in

the pure iron-nickel alloys (up to 6.2 percent

nickel by weight); taenite, corresponding to the

pure 7-phase; and plessite, which may consist

of a microscopic octahedral arrangement of

kamacite and taenite or may be of a granular

nature. If the meteorite contains more than

6 percent nickel, kamacite and taenite bands

appear interlaced in an octahedral pattern,

known as Widmanstatten figures. Occasionally

an iron will have inclusions of iron sulfide, iron

phosphide, diamonds, graphite, and other

minerals. The stony irons are of two kinds,

those in which the iron is distributed as nodules

and those in which veins of iron are intertwined

with siliceous material. The stones consist

largely of siliceous material, through which

are interspersed only flakes or fine veins of iron.

Meteorites have been described in a variety

of ways. X-ray studies have established the

crystal structures of kamacite and taenite

and the crystallography of the Widmanstatten

figures. Neumann bands have been identified

as mechanically twinned regions. Gross compositions and densities of a large number of

meteorites are known. The relative abundances of elements, particularly the rarer ones,

have been found by the use of neutron activation and isotope dilution techniques.

The classical method of fixing the age of a

meteorite is to determine the relative amounts

of helium, uranium, and thorium it contains.

Originally, investigators assumed that all of

the helium was derived from radioactive

disintegration of uranium and thorium. Recently, the existence of helium generated by

81
82 SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS

cosmic rays has been demonstrated and allowed

for in age calculations. An argon-potassium

method, which is limited in application to

stony meteorites, an isotopic lead composition

method, and a strontium-rubidium method are

all in use. The apparent ages of the stony

phase of meteorites determined by each method

are in substantial agreement, approximately

4X109 yr. Iron meteorites are much younger,

3X10* yr, according to the uranium-helium

method; but the tritium measures by Fireman

on recently fallen irons indicate that some

billions of years would have been required for

cosmic rays to produce the observed He8.

These age discrepancies need resolution. Is it

possible, as Urey has suggested, that radioactive He escapes while cosmic-ray He remains

—a consequence of radioactive substances

crystallizing only on grain surface?

Physical metallurgy has contributed significantly, though not extensively, to our

knowledge of the origin and age of meteorites.

The approach here is indirect. Consideration

is given to the combined pressure and temperature conditions, aging, and constraints that

could have produced the observed metallurgical structures. Especially important are

the Widmanstatten figures, the local concentration of nickel within elements of the figures,

and phase transitions indicative of pressure

changes. Research in this area has been exploratory and fragmentary. The need for more

work both on terrestrial alloys and meteorites

is great. The influence of pressure on the phase

diagram of the iron-nickel system is not known

nor have the diffusion coefficients of nickel in

iron been established. The segregation of

elements on a microscale is a completely untouched though most fertile field.

Neumann bands, slickensides, distorted layers, crushed nodules, cleaved surfaces—each

has its story to tell if we can but read it.

Mechanical metallurgy is almost completely

unexplored.

At least 10 authenticated large meteorite

craters exist on the earth. The moon contains

some 2,000 craters believed to be of meteoritic

origin. These craters are important geological

and selenographical features. They are, moreover, positive evidence of encounters between

substantial chunks of fast-moving planetary

and interplanetary material. Study of the

craters gives us much insight into the cosmological role of such physical encounters. It is

only recently that such craters have been

recognized for what they are, and it is still

difficult in most instances to state unambiguously

the velocity and mass of the meteorite that

made the crater. Indeed, it is frequently

difficult to know for certain that the crater

is of meteoritic origin. If we are ever to

know the rate of accretion of matter by the

earth, we must first solve the riddles that

reside in and about the craters. We must

examine the geologic havoc that has been

wrought, we must collect meteoritic debris

from within and without the crater, and we must

conduct high-speed impact tests of our own.
Meteors

By Fred L. Whipple *

In these days of expanded research potentialities, the study of meteors has matured as a

research area and many of the major problems

of yesterday are now solved; for example, the

controversial problem of hyperbolic meteors.

Both the precise photographic measures of

brighter meteors and the radar measures of the

fainter meteors have failed to prove the existence of hyperbolic meteors. If such meteors

exist, they represent a minority below the 1-

percent level. Furthermore, the photographic

and visual meteors are clearly of cometary

origin except for a small fraction not exceeding

10 percent that may be of asteroidal origin.

Most photographic and radio meteors, however, move in direct orbits of relatively low

inclination to the ecliptic; their aphelia tend to

lie beyond the asteroid belt and usually beyond

Jupiter's orbit, but many of the fainter radio

meteors exhibit quite small aphelion distances.

The cometary meteoric mass is proved to consist of extremely fragile, crumbly material that

is probably of very low density. The spectra by

P. M. Millman indicate a composition something like that of stony meteorites, but it is unlikely that any sizable pieces have reached the

ground. Micrometeorites, on the other hand,

may well represent the cometary solids.

Although photographic and radio techniques

have been exploited to an amazing degree in the

past several years, serious gaps remain to be

bridged before we can reach a satisfactory

understanding of meteoric phenomena and of

the origin of these tiny bodies from interplanetary space. I should like to mention briefly

certain of these difficult observational problems

and devote most of my attention to the almost

untouched areas in laboratory and theoretical

work. As long as these areas are unexplored,

our knowledge of meteors, their nature, and

1

 Harvard College Observatory, Cambridge, Mass., and Director,

Astrophyslca) Observatory, Smithsonian Institution.

their origin will remain conspicuously incomplete.

In the area of two-station direct photography

of meteors for velocities, trajectories, deceleration and light curves, the J. B. Baker SuperSchmidt cameras in the Harvard Meteor Program have largely completed the basic observational research in terms of their present-day

capabilities. Analysis is being carried along

rapidly and effectively by L. G. Jacchia, R. E.

McCrosky, and A. F. Cook. I must point out,

however, that no serious effort has yet been

devoted to the use of the television image tube

in optical meteor methods. There is every

reason to believe that the sensitivity can be increased more than an order of magnitude, since

the transient character of the meteoric radiation is highly suited for the present-day photoelectric techniques. Utilization of the image

tube, on the other hand, requires a major effort

since the telescope and the receiver must be

integrated into a new and specialized form.

Such a technique shows promise of reaching to

the ninth magnitude visually for very slow

meteors, and to the sixth apparent magnitude

for the fastest meteors. Applied as in the Harvard Program, the image tube may accomplish

a corresponding improvement (some four magnitudes) in the registration of persistent meteor

trains shortly after the meteors have passed.

Thus, measurement of upper winds as a function

of altitude in the range from 80 to 112 kilometers could be systematic and highly precise.

The train type of observation is important

not only for the understanding of upper atmospheric turbulence but also for solving the basic

problem of meteoric masses. The "coasting

motion" in meteor trains presents to date the

only direct method of measuring meteoric

masses, via their momentum transfer to the

ambient air. Either photographic or imagetube methods should be applied with larger base

83
84 SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS VOL.1

lines than those of the present stations, about

50 rather than 20 miles, to clarify the question

as to the average masses associated with meteors of a given brightness and, from a more

fundamental point of view, the question of the

densities of cometary solids. A single example

leads to a density of only 0.05 gm/cm3!

It has not yet been possible to utilize the

full power of the Baker Super-Schmidt meteor

cameras for the photography of faint meteor

spectra. Relatively coarse transmission replica-gratings with a high concentration of light

in one first order have not yet become available

for use on fast cameras of large aperture.

Coupled with image tubes, such coarse gratings

might lead to remarkable progress in the understanding of the spectra of fainter meteors and

of train and wake phenomena for brighter

meteors and might make it possible to determine the composition of cometary solids.

D. W. R. McKinley at Ottawa, Canada, and

A. C. B. Lovell with his colleagues at Jodrell

Bank, England, have made remarkable strides

in measuring the velocities of faint meteors by

radio techniques. J. G. Davis has now developed a three-station method for measuring

radiants, and, therefore, orbits as well. Only a

major effort utilizing this technique can provide the many data critical to our understanding of the orbits and origins of meteors, as well

as their interaction with the atmosphere. In

particular, it is essential that such a system

be used with a number of receiving stations in

order that observations may be made along the

full length of meteor trails to determine ionization as a function of distance and height, diffusion effects at great altitudes, deceleration in

the atmosphere, precise correction to outer

atmospheric velocities and orbits, heights, winds,

and other important meteoric observables.

Such a program is being undertaken at Harvard

under the technical direction of G. S. Hawkins.

This technique should be carried to its practical limit.

Fortunately, many radio groups are studying

the properties of forward-scatter by meteors

because of its importance in communication.

In the near future, such information needs to

be collated, summarized, and interpreted with

a view to increasing our understanding of the

nature, decay, and diffusion of the electron

clouds produced by meteoroids.

There is every reason to hope that the

planned Earth Satellite Program of the United

States and the International Geophysical Year

of the National Academy of Sciences will

provide direct above-the-atmosphere measures

of small meteoroids in space, so far undetectable

by any of the meteoritic methods. It is

extremely important that such observations

be made in order to clarify the relationships

among comets, meteors, and the zodiacal light

and the effect of corpuscular radiation on small

particles in interplanetary space.

Although sound observational techniques for

the study of meteors are being utilized vigorously, the opposite can be said of the theoretical

and laboratory research which should now be

underway to consolidate these observational

gains. The basic meteoric phenomenon still

remains unexplained in terms of even an

approximate physical theory. This deplorable

situation exists not only because the problem

is difficult, with many facets, but because no

serious effort has been directed towards the

laboratory determinations of the fundamental

data required in such a theory. Desperately

needed are measures of atomic and molecular

cross sections for dissociation, excitation, ionization, attachment, and recombination for meteoritic and atmospheric elements in the energy

range up to 103ev. These quantities are needed for such elements as silicon, iron, nickel,

magnesium, sodium, calcium, manganese, and

others in air, or, better, in pure atmospheres

of nitrogen or oxygen. Such laboratory measures coupled with the remarkable progress of

recent years in shock-tube research and ultra

high-velocity phenomena should make it possible to utilize the observations of meteors to

establish the beginning, at least, of a satisfactory

basic theory for meteors. The needs in the

area of astroballistics are elaborated in a paper

of this volume by R. N. Thomas.

Artificial meteors, either in evacuated ranges

or at moderate altitudes in the atmosphere,

could provide extremely valuable information

regarding the meteoric phenomena under conditions in which the nature and mass of the

meteoric body would be known. Such research, started under the direction of J- S.
NEW HORIZONS IN ASTRONOMY 85

Rinehart at the Naval Ordnance Test Station,

Inyokern, Calif., should be pursued vigorously

to provide the confirmation of theories based

upon the laboratory measures.

In quite a different area of meteoritic studies,

laboratory measures would also be of extreme

value. Corpuscular radiation from the sun

almost certainty plays an important role in the

disintegration of meteoroids in space. We

require laboratory measurements of the sputtering phenomena produced by protons in the

energy range 10*-10*ev on common meteoritic

surfaces—stony minerals as well as nickel-iron.

A knowledge of the rate of disintegration under

these circumstances would greatly increase our

understanding of the cometary meteoroid.

Once such a body has been ejected from a

comet, it is subjected to corpuscular radiation

which causes it to spiral inward towards the

sun and also to disintegrate. There seems

little doubt that cometary particles are the

major source of scattering and diffraction of

sunlight in the solar Fraunhofer corona, in the

zodiacal light, and in the Gegenschein. We

cannot integrate our information concerning

comets, meteors, the corona, the zodiacal light,

and the Gegenschein until we know the effect

of corpuscular radiation on small meteoritic

particles.

Similar laboratory measures of atmospheric

components impinging on meteoritic material

with energies in the range 10-103ev would

provide a sound foundation for an understanding of the interaction of extremely small

meteoroids with the upper atmosphere. If the

sputtering effects produced under these conditions are sufficiently small, then "micrometeorites" can indeed pass through the upper

atmosphere without destruction; by black-body

radiation at temperatures below the melting

point they can radiate away the heat derived

from atmospheric encounter. Such sputtering

information would be of extreme value in

interpreting the observations of meteoritic dust

in the atmosphere and in deep sea oozes. It

would also help clarify the question of whether

micrometeorites provide condensation nuclei to

initiate heavy rainfall at extremely great

altitudes.

In the laboratory-ballistic area of measurement, particularly needed are the values of such

3S1O14 — 30 7

physical quantities as the heat transfer coefficient and the luminous efficiency. These

quantities are, respectively, the percentage of

the available translational energy of the medium that is transferred to the surface of the

meteoric body and the percentage energy of

the ablated material that is transferred into

observable radiation. A corresponding quantity is the ionization efficiency, measuring the

number of electrons freed in the meteoric

process by each atom of the meteoroid. These

quantities are, naturally, functions of the

velocity and very likely are functions of the air

density and composition of the meteoroid.

The physics of the meteoric process involves

the integration of physical data, physical

phenomena, and meteoric phenomena into a

coherent whole. The development of such a

theory has obviously been the goal of much

research. In one area, there is hope that such a

theory can be developed without so much subsidiary information, viz, the theory of the persistent meteor train- The evidence points

strongly to the conclusion that meteoric energy

is transferred rapidly to some retentive agency

in the atmosphere, such as active nitrogen, and

is transformed slowly into radiation after the

meteoric phenomenon is ended. Such a theory

will be closely integrated with the problems of

the physics of the upper atmosphere but does

not seem to be hopeless in terms of our knowledge at the moment.

A number of theoretical problems fall in

the realm of celestial mechanics. The problem

of the "Jupiter barrier" requires much more

attention. The Jupiter barrier arises from the

perturbational disturbances by a large planet

which tend to prevent a particle spiraling in

towards the sun from reducing its aphelion

distance below the perihelion distance of the

planet. Perturbations by close approaches

to the planet necessarily throw the particle

into an orbit that will return near the point of

encounter, and thus an orbit with aphelion

beyond the planet's orbit. Many of the

perturbational problems of interactions between

the planets and small particles of the solar

system take on a stochastic character and, in

fact, are soluble by modern stochastic theory.

An interesting problem combining stochastic

theory and celestial mechanics concerns E.
SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS VOL.1

Opik's suggestion that perturbations of asteroidal material by Mars are responsible for the

existence today of asteroidal material crossing

the earth's orbit. Here we require a more

comprehensive theory of the perturbative

effects of a smaller planet upon microparticles

in the solar system.

This enumeration of problems in the field

of meteors is far from complete; in particular

it omits a number of important questions in

the category of correlation-interpretationtheory. These questions involve the relationship between the orbits of comets and meteors,

the distribution of particle orbits to produce the

observed zodiacal light, the nature of the Gegenschein, the development of meteor streams,

identification criteria for meteor streams of

great diffusion or low population, and a number

of problems relating to the orbits of comets,

meteors, meteorites, and asteroids.

The author hopes that this report, incomplete

though it is, will at least tend to eliminate

complacency with regard to the present completeness of meteoric research.

References

KAISBH, T. R., BD.

1955. Meteors. Pergamon Press, Ltd., London.

LOTBLL, A. C. B.

1954. Meteor astronomy. Oxford at the Clarendon Press, New York.

POBTEB, J. G.

1952. Comets and meteor streams. John Wiley

& Sons, New York.

RlNEHABT, J. 8.

1954. Behavior of metals under impulsive loads.

American Society for Metals, Cleveland,

Ohio.
Minor Planets

By PaulHerget1

The techniques and methods of observing the

minor planets have not improved substantially

for more than half a century. Within the last

decade, however, the development of electronic

computing machines has increased the possibilities for computational work by a factor of

nearly a thousand. Advances in technology

have improved photographic emulsions so that

we can observe planets much fainter than those

observed formerly, but there is no immediate

prospect of improving methods of blinking,

measuring, and reductions based on stellar comparison positions from star catalogs. Minor

planets will continue to be discovered more or

less at random, and each one must be followed

almost individually for several months in order

to secure its orbit and to predict the following

opposition.

One of the most difficult problems for the

Minor Planet Center occurs when the observer

assigns an erroneous identification to the object

whose position is determined. As fainter and

fainter objects are observed, the likelihood of

making such misidentifications increases greatly

and can be avoided only by securing a confirming observation from two to four weeks later. In

very doubtful cases even a third observation is

needed so that an independent orbit can be

computed, if necessary. For this same reason,

it is now increasingly important to provide

much more accurate ephemeris and perturbation

computations than in the past.

In 1950 Kuiper launched a large-scale observing project in which the minor planet zone

of the sky, about 23° on each side of the ecliptic,

was observed in the neighborhood of opposition

continuously for a period of 20 months. The

main objective of this program was to gather

fairly homogeneous data on the magnitudes of

1

 Department of Astronomy, University of Cincinnati, Cincinnati,

Ohio.

the minor planets and to provide a basis for investigating theories of their origin. A very

important concomitant result was a nearly complete survey of all minor planet positions, which

proved to be a very valuable standard of reference whenever doubtful cases occurred in the

improvement of orbits. If the computations

were based on actual observations it should be

possible to locate the doubtful planet on the

survey plates from its computed position and,

thereby, provide still more observational data,

unless the planet was below the magnitude limit

of the plate. On the other hand, if the supposed

observations were spurious it would be impossible to locate an object at the computed position on the survey plates. The value of such

a survey is so great that it should be repeated at

least once every decade. For the ordinary minor planet orbits, with sufficient observations in

the past and with good perturbation computations, this should be sufficient to assure accurate ephemerides; and it may have the effect of

displacing at least half of the ordinary minor

planet observing that now goes on.

There is room for a special type of observing

program which could be undertaken by someone

who has a sufficiently large telescope, who likes

to sleep through the night, and who is willing to

work low on the horizon. The purpose would

be to increase the length of the observed arc of

newly discovered minor planets, both at the end

of the first opposition in which they are discovered and at the beginning of the next following opposition. A moderately large telescope,

something like a 24-inch reflector, would be

needed because the magnitudes will be unusually faint because of the atmospheric extinction

at low altitudes and the disadvantageous phase

relationship nearer to superior conjunction.

Also, many of the newly discovered faint planets

will be in the region of perihelion during the first

87
88 SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS

opposition and, due to their eccentricity, will be

at a greater heliocentric distance in the second

opposition.

Recent work at the Naval Observatory and

by Kuiper has disclosed new problems related

to the physical constitution and history of the

minor planets. Their solution requires concurrent photoelectric observations of individual

minor planets, including changes in color,

magnitude, and polarization, and extending

over a wide range of phase angles. Some

evidence already exists that two classes of

minor planets can be recognized from polarization observations. Laboratory work of a

similar nature with known materials and conditions should offer some clues to these problems. Theoretical studies are also desirable.

The Minor Planet Center has experimented

with all available types of computing machines

in recent years, and it has now concluded that

the perturbations of all ordinary minor planets

may be most readily handled either by the

method of variation of vectorial constants or

by Hansen's method. The computations required are no longer a formidable problem.

Similarly, the problems of preliminary orbits

and differential corrections have been reduced

to a matter of minutes. The same results

could easily be attained for all comet

computations.

The routine aspects of minor planet computations are so well in hand that it is now possible

to turn to new and more interesting problems,

the solution of many of which will be aided by

the experience gained in the work described

above. An accurately computed dynamical

orbit of Pallas from the time of its discovery

up to the present will permit an independent

and highly reliable determination of the value

of the constant of precession. G. W. Hill

prepared a list of 13 planets that are especially

suitable for determining the mass of Jupiter,

provided the observations and accurate perturbations extend over a period of about a

century; work on this problem has begun.

Another program, already well underway, will

use the first four minor planets for determining

the position of the equinox instead of observations of the sun and the major planets. This

project also requires accurate computations

of their orbits. More penetrating researches

may now be made upon the families of minor

planets originally recognized by Hirayama.

Electronic machines have not yet been used

as effectively for the calculation of general

perturbations as for that of special perturbations,

but the problem has been attacked and will

undoubtedly result in progress in the near

future. The availability of electronic computing machines now makes the vistas in all

these fields more attractive than at any

previous time in history.
Planets, Satellites, and Comets

By Gerard P. Kuiper *

The planets and the satellites must be

studied from distances (—1 astronomical unit)

that are intermediate between that of the stars

(>105 a. u.) and that of the bulk of the earth

(< 10~6 a. u.). Likewise, some of the problems

they present are intermediate—related both to

astrophysics and geophysics—but this situation

is rather recent.

The following paragraphs, part of a preface

to a book now in preparation, may serve as a

backdrop on which to view some of the 20th

century problems of planetary astronomy.

Until the introduction of photography with large

telescopes and the birth of astrophysics, both toward

the close of the 19th century, planetary astronomy—

with some exceptions—was astronomy. The system of

the "fixed" stars was of profound interest philosophically; but to the practicing astronomer it was largely a

convenient frame to which the complex motions of the

planets could be referred. The great discoveries in

astronomy were made in efforts to interpret the planetary motions. Copernicus, Kepler, Newton, Euler,

Lagrange, Laplace, Gauss, and Poincar6, to name but

the greatest, created the concept of modern natural

science while studying the planets.

The phenomenal growth of astrophysics, and the

exciting explorations of the galaxy and the observable

universe, led almost to an abandonment of planetary

studies. The number of astronomers has always been

small and for better or for worse has not followed the

enormous increases of physicists and chemists. Celestial mechanics began to look like a nearly finished

discipline in which further progress was hopelessly

difficult. Physical observations of planetary surfaces,

particularly of Mars, led to controversies and speculations. More and more this branch of planetary work,

including the study of the moon, became the topic

par excellence of amateurs—who did remarkably well

with it. The Memoirs of the British Astronomical

Association became the chief record for the development of planetary surface markings. Astronomers with

large telescopes were so occupied with the engaging

problems of stars, nebulae, binaries, clusters, the

galaxy and the universe, that astronomy almost entirely became the science of the stars.

1

 Yerkes Observatory, Williams Bay, Wis.

In some respects the science of planetary astronomy has greatly changed from its distinguished predecessor. Eight- or ten-decimal accuracy and a formidable theoretical apparatus

have often given way to one-decimal accuracy

and order-of-magnitude arguments. Observing programs, which occupied the major observatories for decades, have been replaced by

isolated observations by single astronomers who

build new apparatus and use it once or twice

under optimum conditions. "Astronomical accuracy" is a compliment born in an age different

from ours.

Nevertheless, there has been no breakdown

of standards. What has happened is that a

new branch of science was born, an offshoot

of young and vigorous astrophysics, which was

added to classical planetary astronomy but was

not immediately integrated with it. And so

it happened that 10-decimal accuracy and

1-decimal accuracy existed side by side, in a

field now occupied by scientists of very different

interests and training. The study of planetary

motions and the great problem of the stability

of the solar system rose slowly and laboriously

on foundations laid by the great masters.

New computational techniques proved increasingly valuable in the solution of special problems. Dr. Clemence reviews this basic part of

planetary astronomy elsewhere in this series.

The program of this 20th century addition

consists of a great variety of things that at

first sight may seem to lack coherence. A

geophysicist, considering the task of a planetary astronomer, might well be skeptical.

He could reason that, since the objects of study

are some 10s times more distant than the earth,

the obtainable information might be 1010

times smaller, and furthermore, that the absence of material contact or near-contact with

the planets deprives the astronomer of several

sources of information: seismic studies, heat

80
90 SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS VOL.1

flux measurements from the interior, gravity

surveys, magnetic surveys, study of magnetic

variations, volcanism, chemical and crystallographic analyses of the crust, oceanographic

studies, detailed atmospheric studies, etc.

One might wonder what there is left for the

astronomer to do.

Classical or gravitational astronomy, however, has put the planets on a coequal basis

with the earth in some aspects. The mass of

the earth in absolute units is known no better

than the masses of Venus, Mars, Jupiter,

Saturn, Uranus, and the Moon, with Mercury

and Neptune closely following. Nor do we

know the planetary orbits with an accuracy

greatly different from that of the earth's

orbit. Hence, the radiation received from the

sun is known equally well for all. The albedos

of the planets can be determined with considerable precision; hence the equilibrium

temperatures can be computed quite well,

even if with somewhat higher precision for some

of the planets than for the earth, whose albedo

is difficult to determine accurately. The

periods of rotation of the planets can be found

(Venus is still an exception), often with considerable precision, as well as the planetary

obliquities. With the radiation intensity and

the obliquity known, theoretical climatology

may be applied to the planets provided something can be said about planetary composition.

This, too, is within the range of possibilities.

It is immaterial whether the source is 109

miles away or just in front of the slit of the

spectrograph, provided enough light is available to make a spectrogram of the desired

dispersion during, say, one night's exposure.

There are, however, serious handicaps, not

resulting from distance but imposed by Nature

itself. In the first place, some of the most

abundant atmospheric constituents are homonuclear molecules (H2, N2, O2), and because of

their symmetry the vibration-rotation spectra

are forbidden. Since the electronic absorption

spectra H2 and N2 are in the ultraviolet part of

the spectrum that is cut off by our atmosphere

(X<0.3M), we cannot discover these gases on

other planets unless such enormous quantities

are present that pressure effects destroy the

symmetry of the molecules and make their

infrared spectrum faintly visible (as on Uranus

and Neptune for H2). Nitrogen in our atmosphere is a good example of invisibility; even at

sunset, when the sun's rays pass through 320 km

atm., four-fifths of which is N2, it does not add

telluric lines to the solar spectrum, even at high

dispersion. A second handicap affects the inert

or noble gases, He, Ne, A, Kr, Xe. They are

of great interest for the study of the origin of

the planets because they will not have combined

with other elements during geologic time and

must therefore always have been part of the

planetary atmospheres. (This statement does

not apply, of course, to radiogenic He4, He3,

and A40.) Because these gases range in atomic

weight from 4 to 130, we can use them—in the

case of the earth, where they can be observed—

to determine empirically how the fraction of

atmospheric constituents lost by evaporation

during the process of planet formation depends

on molecular weight. Since the original composition of the protoplanets, from which the

planets condensed, was very probably the cosmic composition (essentially the composition of

the solar atmosphere), and since the abundance

of heavy elements, like Fe, Mg, Si, etc., may be

estimated, we may compare the terrestrial

abundance table with the cosmic abundance

table to obtain the desired ratio, assuming that

no appreciable fraction of the heavy elements

was lost. Now it would be exciting if one

could apply this approach to other planets.

But the very property which makes the noble

gases so interesting in tracing the early history

of the earth (i. e., their stability in the atomic

state due to the completeness of their outermost electronic shell) makes them at present

undiscoverable on other planets. Such high

energies are needed to excite these atoms that

their absorption spectra are far in the inaccessible ultraviolet.

It is clear, then, that two important advances

may have to await the development of an

extraterrestrial observing station. If planetary

spectra could be taken from 200- or 300-km

elevation, entirely new and vital information

could be gotten on the composition of planetary atmospheres. One would then get an

undistorted picture of atmospheric composition.

The handicaps imposed by Nature would have

been overcome.

Some of this much-desired information can
NEW HORIZONS IN ASTRONOMY 91

probably be obtained with a large telescope at

a ground observatory. By introducing compensation methods, Hiltner has succeeded in

measuring polarizations down to a few units in

the fourth place. It should be possible, in

principle, to do the same in spectrophotometry.

This would allow one to observe the Raman

spectra of H3 in the Jovian planets and to make

a rough determination of the H2/CH4 ratio.

Conceivably even the He abundance might

thus be found. Similarly for the terrestrial

planets, N2 and possibly other gases might

become observable, at least if they are more

abundant than the strongly scattering CO2. It

even seems possible, in principle, that the solar

lines Mg II at XX2795.5 and 2802.7 A, including

their variable emission components, could be

observed regularly from the ground from the

Raman lines in the spectrum of Jupiter, at

XX3163 and 3172 A. The problem is probably

largely a question of instrumental development.

Related to the problem of atmospheric composition are at least four others capable of

empirical pursuit. They are the problems of

atmospheric clouds, the planetary albedo and

its wavelength dependence, planetary polarization as a function of the phase angle and as a

function of wavelength, and the spectrum of

the infrared emission of the planet to space,

which at least for the terrestrial planets must

balance the absorbed solar radiation. For the

Jovian planets a finite fraction of the radiation

emitted probably still derives from gravitational

contraction and from cooling off of the planetary

interior.

The atmospheric clouds on Jupiter and Saturn

are almost certainly atmospheric condensation

products, as is true for most of the clouds and

haze on Mars. Occasionally Mars has yellow

dust clouds also, while the nature of the cloud

cover on Venus is still in doubt. Uranus and

Neptune probably have no visible clouds, although markings have been reported on Neptune. The reality of these markings is not confinned by observation with large telescopes, or

by the model required for the atmospheres of

Uranus and Neptune to account for the enormous quantities of CH4 and H3 observed. Saturn has a general haze cover which forms the

bottom of the visible atmosphere. It usually is

divided into a number of distinct belts, parallel

to the equator, of slightly different reflectivity

and, sometimes, color. In addition, separate

clouds are occasionally seen, perhaps one per

decade on the average; they must, of course, be

huge to be visible, with diameters about as large

as the entire earth. They seem to be caused by

major eruptions from the interior and, because

of the very pronounced differential rotation

within the planetary atmosphere, they are soon

stretched out into belts parallel to the equator,

merging with the belts already present. It is

quite interesting how the belt system of Saturn

has changed with time. Jupiter, likewise, shows

major upheavals in its cloud cover, every decade

or so, not at all related to its position in its eccentric orbit, and therefore not caused by solar

heating. These two planets apparently do not

have a sun-induced meteorology, as do terrestrial planets, but one largely governed by released internal heat. It has been pointed out

by meteorologists that internal heating is much

more efficient in stirring an atmosphere than

solar heating from the outside.

Fascinating and unsolved problems are posed

by these observations. Do these occasional

bursts come from the deep interior? Then how

do they get through the mantle of solid hydrogen usually assumed to surround the denser

core ? Or is contraction of the mantle itself the

cause ? What is the explanation of the brightly

colored clouds—brick red, chocolate brown,

blue, amber, white, etc.—observed especially

during these upheavals? Is the Red Spot a

cloud cap over a floating island, and if so, what

floats in what? Or is it merely a long-lived

vortex in the Jupiter envelope? (It is not attached to the core because its rotational speed

is variable.) These questions, in turn, are related to the unsolved problem of the depth of

the gaseous envelope and the hydrodynamics of

the differential rotation on both Jupiter and

Saturn. Of the greatest interest and promise is

the recent discovery of a whitish cloud on Jupiter emitting radio bursts at irregular intervals; the power of the bursts measured at 22 Me

is 10* times that of a terrestrial lightning

discharge, but the spectrum is quite different.

The cloud problem for Mars is simple by

comparison, and has been largely clarified by a

combination of the brightness at different wavelengths, polarization measures, and studies of
92 SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS VOL. 1

general behavior. By contrast, the identity of

the Venus cover is not at all obvious. Different

suggestions have been made: water droplets,

dust, and salt particles, but none of these fit

all the observational data on polarization, in

visual and infrared light, on the brightness

variations with phase, on the planet's color,

and on the appearance of the cloud belts. In

conversations with Professors W. E. Grath of

Bonn and P. Harteck of Rensselaer Polytechnic

Institute, the possibility was discussed that the

particles are polymerized C3O2, formed as a

result of the impact of solar ultraviolet radiations on an atmosphere largely composed of

CO2. This possibility will be checked by laboratory experiments. Here is an exciting problem that, once it is solved, will also deepen our

understanding of the history of the earth, so

similar to Venus in dimensions and mass.

A curious problem is that of the planetary

radiation to space. For a planet or a satellite

without an atmosphere, there is no difficulty.

Black-body radiation will be a good approximation for such an object, covered as it probably

is with fine powder or debris after its long

exposure to a vacuum, to solar radiation varying

with the rotation, to cosmic rays, and to interplanetary debris. But how does a planet like

Uranus or Neptune establish its temperature

balance with the absorbed solar radiation?

What atmospheric constituent radiates effectively between 20M and 100/*, where the maximum of the emission should occur? With

H2 a homonuclear molecule and CH4 a symmetrical molecule, and NH3 and H2O frozen

out, a situation arises which is not unlike that

in the upper layers of the earth's atmosphere.

There, too, appreciable quantities of (in this

case ultraviolet) solar radiation are absorbed,

but the atoms cannot dispose of this energy

by radiating in extensive infrared bands. The

result is that the atmosphere has a hot upper

fringe at nearly 10 times the mean equilibriumradiation temperature of the earth ( ^ 250° K).

The thermal radiations in the radio-frequency

range will be of special interest since they will

presumably arise from the solid surfaces, and

therefore give a measure of the depths of the

atmospheres. Thermal radio emission of Mars

will determine the mean daily temperature and,

by implication, the surface temperature at

night.

Enough has probably been said to illustrate

the wide variety of planetary problems, many

of them unsolved. One could mention others:

precise diameter determinations and resulting

planetary densities; their interpretation in

terms of models based on the physics of matter

under very high compression; determination

of the moments of inertia and the resulting

refinements of the planetary models; the related

problem of planetary composition; and relation

of the interior to the envelope for the Jovian

planets. In a separate class is the problem of

the dark Martian features and their variations,

and the question as to whether they are organic

or inorganic.

Satellites present a host of problems of their

own. Needless to say, the moon is foremost

among them, but they are all different and

interesting. Titan, with its orange color and

its atmosphere of methane; Triton, with its

slightly yellowish color and its probable atmosphere, of unidentified composition; Io,

with its bright orange color, but lack of atmosphere; the snow satellites of Saturn; the

greatly different mean densities and, hence,

compositions of the satellites; the great variety

of albedos, from near unity for the inner

satellites of Saturn to 0.07 for the moon; the

differences in albedo sometimes found for

different hemispheres of the same body, exceptionally large for Iapetus; the spots on the

Jupiter satellites and on the moon; etc. The

Rings of Saturn are unique and of unrivaled

beauty in a large telescope; their snowy composition seems reasonably well established.

One can understand why Jupiter does not have

such a ring. The high density of Io indicates

that the surroundings of Jupiter were quite

hot at the time Io formed; and, even if some

snow later condensed within the Roche limit

of Jupiter, it would almost certainly have

evaporated subsequently by solar radiation.

Saturn's ring appears to be just safe from

such destruction. But why don't Uranus and

Neptune have rings? That is very puzzling

and further work must be done to lower the

thresholds set by present observations.

Comets are much stranger objects than

planets or satellites. Their ephemeral nature
NEW HORIZONS IN ASTRONOMY 93

shows that they "do not belong" near the

earth; that they were not formed there. The

exposure to strong solar heat and corpuscular

radiation causes the comets to disintegrate

rapidly, and show heads and long tails that are

among the weirdest structures seen in Nature.

No wonder pamphlets were written in past ages

on whether the appearance of a comet might

mean the death of a king. Comets are unadjusted to their surroundings, as we observe

them near the earth; for this very reason they

demonstrate, by their violent reactions, certain

properties of these surroundings, as the old

and adjusted planets and satellites no longer do.

Comets have thus become of the greatest

interest in the study of the properties of solar

radiations and interplanetary magnetic fields,

as well as for their own sake. The unpredictability of the occurrence of a "new" comet and

the irregularity of the stimulating solar beams

add to the excitement when a bright comet

appears, which has not been very often in

recent years. A good bright comet, like that

of 1882, would be a major scientific event.

Problems of evolution and origin are as compelling in solar-system studies as they are in

studies of stars and galaxies. It seems as if

some of these are being solved. On the other

hand, a good look at the problems in geophysics

will temper any excess optimism. Much of

what we would like to know about planets,

satellites, and comets is hidden from our view;

and even if we could see it, explanations would

not follow automatically, as geophysics shows.

Clearly, one must be satisfied with limited

objectives and with explanations and theories

of a general kind. But because of the many

repercussions these problems have on geophysics

and related sciences, any advance will also have

a general interest. The limits imposed by

telescopic equipment and other practical considerations have by no means been reached.

381014—56-

Airglow and Aurora

By A. B. Meinelx

The fact that the background of the moonless

night sky is not entirely dark has been known

for centuries, while references to the aurorae

extend back into the dim recesses of Nordic

antiquity. Yet only since the turn of the 20th

century have these phenomena been the subject

of systematic study. The presence of an

emission spectrum for the aurora was recognized

as early as 1867 when Angstrom observed and

measured the wave length of the green auroral

line. The presence of an emission spectrum

from the light of the night sky was noted as

early as 1895 by W. W. Campbell, who used a

new visual spectrograph at Lick Observatory

many years before the spectrum was rediscovered photographically by V. M. Slipher at

Lowell Observatory. The light of the night sky

as a detached phenomenon of physical interest

was recognized by the pioneering investigations

of Fabry in 1910, V. M. Slipher in 1919, Lord

Kayleigh in 1920, and Dufay in 1923. For a

historical account of early night sky and auroral

research, the reader is referred to the articles by

Dufay (1928), Fabry, Dufay, and Cojan (1934),

Dejardin (1936), Elvey (1942), Dufay (1938),

and Swings (1949).

Many young scientists may wonder why

airglow and aurorae are included in a report on

new horizons in astronomy. Historically, both

phenomena are noteworthy because of their

astronomical consequences, although both arise

solely within the gaseous envelope of the earth.

To most astronomers the light of the night sky

is a nuisance to be tolerated by necessity, since

little can be done about it. However, in spite

of the unpleasant problems they make for the

observer, the light of the night sky and the

aurorae are phenomena of wide astronomical

and astrophysical interest.

The term "light of the night sky" is used to

'National Astronomical Observatory, and Yerkes Observatory,

Williams Bay, Wls.

describe a complex flux of radiation composed

of both terrestrial and extraterrestrial sources.

The extraterrestrial component arises from:

(a) interplanetary scattering of sunlight (zodiacal light), (b) interstellar scattering and emission excited by stars, and (c) the integrated

radiation of individual unresolved stars. The

atmospheric component of the light of the night

sky arises from: (a) scattering of starlight and

extraterrestrial radiation in our atmosphere,

(b) atomic emissions in the form of monochromatic emission lines, (c) molecular emissions in the form of emission bands and possibly

continua, and (d) Cerenkov radiation and

N2

+

 emission from cosmic rays. Assigning to

each source a percentage of the total light of the

night sky can only be done with uncertainty since

it is difficult to separate clearly some effects.

In the visible region the ratio of extraterrestrial

to terrestrial must be approximately unity.

Upon certain occasions aurorae are superimposed on the night sky. The auroral radiations are caused by two basic mechanisms:

(a) heavy-particle excited atomic and molecular

emissions (in the homogeneous arc phase), and

(b) electrical-discharge excited atomic and

molecular emissions (in the rayed-structure

phase). However, most of the light may be

produced by secondary electrons in both cases.

To describe the appearance of the light of

the night sky on the spectrum of a faint astronomical object, such as a distant nebula, it is

best to portray each spectral region separately.

In the region sensitive to the 1030^-0 emulsion,

the XX4800-3100 spectrum is characterized by

many bandlike features superimposed upon a

continuum which is strong in the blue and

rapidly disappears shorter than X3900. As a

consequence of this spectrum it is difficult to

detect weak astronomical features such as the

broad, but shallow, H and K lines in the spectrum of faint galaxies.

05
96 SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS VOL. 1

The region <X3500 is composed of groupings

of bands of increasing intensity and contrast

down to the atmospheric cutoff. These bands

appear distinct and shaded toward the red, in

contrast to the diffuse Xi and X2 bands in the

X4500 region. No emission lines of atomic

origin occur in this region except during auroral

activity.

The visible region of the airglow spectrum,

XX6500-4200, is dominated by four sharp emission lines with relatively weak bands and a

continuum near the blue end. The lines, in

order of intensity, are the famous auroral line

at X5577, two red nebular lines at X6300-

6364 (all three due to forbidden 01 emission),

and the sodium doublet at X5893. The latter

is generally so weak that it is only recorded on

low-dispersion spectrograms and as a single

line. The other features are weak OH bands

and an unidentified continuum that must be

heavily contaminated by starlight, etc.

Whether this "continuum" is real and of

atmospheric origin is still unknown. On spectrograms sensitive to the deep red (i. e., 103a-F

emulsions) the OH bands near X6560 and X6800

become prominent features.

Spectrograms of astronomical sources in the

visible often show additional lines that could

puzzle the new observer. These lines are due

to scattering by the atmosphere of neon and

mercury radiations from adjacent cities. The

relative intensities generally appear anomalous,

especially for neon in the red because of the

"sunset" reddening over long air paths due to

rapid changes in the absorption coefficient.

In the infrared, the airglow becomes intense

as a result of a strong band system due to the

OH radical. The most intense bands of this

system are 4.5/*, but the bands at X10,000 and

in the I-N emulsion photographic region

(XX8900-6500) are strong enough to cause difficulty in some astronomical observations.

The auroral spectrum shows the same atomic

line features as the airglow, with the addition

of a rich assortment of strong molecular bands

in the red and infrared and violet regions.

The absolute intensity level can reach a factor

of 10* times that of the airglow. Weak aurorae

often make themselves evident on astronomical

spectrograms only by enhancement of the X6300

[01] lines and the appearance of the X3914 N2

+

band, which appears as a line owing to the

occurrence of the band head in the red branch.

Both the airglow and aurorae are phenomena

that are potentially important to the understanding of the physical and chemical properties

of the upper atmosphere. Our atmosphere

provides us with a nearby astrophysical laboratory where we can observe a gaseous atmosphere

over a wide pressure range under the influence

of electromagnetic and particle radiation from

a star. Unfortunately, the interpretation of

these phenomena has proved so complex that

we are still in doubt with regard to many of the

questions that investigators have struggled

with. To begin with, both phenomena, airglow and aurora, are difficult to study with

resolution sufficient for detailed analysis. The

airglow is very weak, necessitating the use of

the fastest possible spectrographs. Even so, a

single exposure sufficient to resolve molecular

bands may require all the dark hours for several

months. Although the aurora is much brighter

(up to 104), it is transient in both time and

space, sometimes lasting for the entire night

but more frequently only a few minutes. Quite

often the lifetimes of interesting features are

measured in tenths of a second. Newly developed photoelectric techniques may offer hope

in attacking such problems.

For some years, the major airglow problem

has been the identification of the emissions.

This problem is approaching what is thought to

be a solution, although X, and X2 bands of

Lord Rayleigh in the blue are still completely

unidentified since no one has yet resolved their

rotation structure, if they possess any. Among

other interesting problems are the heights of the

emissions of the night glow and the excitation

mechanisms of both.

In the case of heights, very ingenious instruments have been devised to measure the patchiness and motions of the airglow patches. Estimates of the heights based upon the dependence

of brightness on zenith distance and made by

averaging many nights of observations with

careful accounting for absorption and scattering

have been disappointingly at variance. The

reason may lie in the fundamental assumptions

of a statistically uniform, thin, emitting layer.

In the case of excitation mechanisms, some

emissions have been thoroughly explained, while
NEW HORIZONS IN ASTRONOMY 97

others present unanswered questions. The

aurora is in particular interesting since the spectrograph has demonstrated the influx of energetic protons, presumed to be of solar origin.

The precise role of the protons in direct excitation of the atmospheric gases taxes the present

capacity of quantum mechanics, since accurate

wave-function descriptions of the molecular energy states are still meager. Laboratory experiments stimulate the imagination, and when

such experiments have been carefully devised

the results have not been disappointing.

In the immediate future, the program for the

International Geophysical Year will attempt to

gather data on a world-wide scale with the best

means uncovered in individual research studies.

The reason is that one-point observations, even

in great detail, leave phases of the problem presented by the aurora unanswerable. The synoptic observations that will be obtained should

open a new door to the interpretation of solarterrestrial relations; however, as in the past,

the unlocking of these secrets will be a challenging but rewarding task for names not yet on the

pages of the record.

General references

CHAMBERLAIN, J. W., AND MBINEL, A. B.

1954. In Kuiper, ed., The earth as a planet. University of Chicago Press.

DEJARDIN, G.

1936. Rev. Mod. Phys., vol. 8, p. 1.

DUFAT, J.

1928. Bull. Obs. Lyon, vol. 10, No. 9, p. 81.

1938. Bull. Obs. Lyon, vol. 3, No. 16.

ELVEY, C. T.

1942. Rev. Mod. Phys., vol. 14, p. 140.

FABBT, C, DUFAT, J., AND COJAN, J.

1934. Etude de la lumidre du ciel nocturne.

Revue d'Optique, Paris.

MBINEL, A. B.

1951. Phys. Soc. London Rep. Prog. Phys., vol.

14, p. 121.

SWINGS, P.

1949. In Kuiper, ed., The atmospheres of the

earth and planets, ed. 1. University of

Chicago Press.

SWINGS, P., AND MEINEL, A. B.

1951. In Kuiper, ed., The atmospheres of the

earth and planets, ed. 2. University of

Chicago Press.

Solar-Terrestrial Relationships:

Weather and Communications

By Walter Orr Roberts*

Relatively few of the many branches of

astronomy can be said to have applications of

direct practical importance. The field of solarterrestrial relationships is an exception. Two

decades of practice have shown the radio

operator that solar variations can make the

difference between success and failure in getting

a message through. Today, we may look for

improvements in practical weather forecasting

because of new discoveries made on the common

border between meteorology and astronomy.

Just how far away these advances are will

depend, at least in part, on the magnitude and

the quality of the effort made in this field.

The effects of the sun on radio communications are well known, even if poorly understood.

The influence of solar activity on weather is

less apparent, and the reality and nature of the

effects are still somewhat controversial. Enough

clues exist, however, to convince a substantial

group of research workers that changes in the

sun's shortwave emanations and in its corpuscular radiation are reflected in worldwide

weather patterns, and that practical gains in

both long- and short-term rainfall forecasting

and temperature forecasting for large areas of

the earth can be expected when these connections are better understood. The fact that the

sun exhibits rather stable long-term trends

improves the prospects for season-in-advance

weather forecasts based on variations in solar

activity.

The sun's effects on weather obviously must

be on a large scale. Several factors in the socalled "general circulation" of the earth's

atmosphere give evidence of such solar influences. The results of these large-scale

changes, in terms of surface rainfall and

• High Altitude Observatory, Boulder, Cote.

temperature, depend very much on local elements such as topography and latitude; a

worldwide circulation change that produces

rain at one spot may produce drought in

another. The ties of sun to local weather are

therefore most complex.

The mechanism responsible for responses of

weather to solar activity clearly does not involve substantial changes of the sun's total

energy; rather, it involves changes in selected

spectral ranges where only a relatively small

portion of the sun's total energy output is

concentrated. It is also clear that the most

pronounced weather responses are likely to be

found at the upper levels of the earth's atmosphere, where the sun's variable energy is

principally absorbed. However, meteorological observations at such levels are still rather

sporadic. Improvements in all types of weather

forecasting are likely to come from research on

the nature and causes of persistence and sudden

changes in the dynamics of the general circulation of the atmosphere. The motions of the

upper levels and the possible nature of the

mechanisms that start disruptions in the patterns of momentum and mass transport in the

earth's atmosphere require particular attention.

Dynamic meteorology is still in its infancy, but

the solar astrophysicist has much to contribute.

Ten years ago there seemed no real prospect

that major advances in long-term weather

forecasting would result from solar observations; today they promise advances of incalculably vast practical significance.

Without much exaggeration it can be stated

that modern solar-terrestrial research began

when R. C. Carrington of the Royal Observatory in England made his classic observation of

a great solar flare visible in integrated light at
100 SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS VOL. 1

11 a. m. on September 1, 1859. At the instant

of the flare, a fluctuation occurred in the

strength of the earth's magnetic field. On

the 2 days following, a violent magnetic storm

was recorded by magnetometers at the Kcw

Observatory. Balfour Stewart, the director

at Kew, concluded that the terrestrial magnetic

disturbances were related to the flare; however,

the idea was so unbelievable that most research

workers dismissed the sequence of events as a

coincidence.

The connection between flares and magnetic

storms is now amply demonstrated, although

the problem remains of explaining just how a

solar flare produces these effects and the even

more spectacular phenomena in the aurora

and ionosphere. In the years since Carrington's great flare, a whole new science has

developed to solve the problems common to

solar physics and to terrestrial atmospheric

physics.

Both sides of this joint attack have greatly

enriched our observational knowledge. The

spectroheliograph, the coronagraph, and the

solar radio telescope have been the astronomical

milestones. Geophysically, the night sky spectrophotometers and ionospheric sounding devices have competed in importance with earthbased, airborne, and rocket-borne cosmic-ray

counters, with photometers, and with magnetometers operated with ever-increasing regularity from an ever-increasing diversity of stations.

Great advances, particularly in meteorology,

have come from the simple fact that we now

can prepare daily or twice daily hemispherewide charts of airflow and temperature distribution at a wide range of altitudes from

which large-scale dynamical studies can be

made.

In every facet of terrestrial atmospheric

physics, tantalizing connections between sun

and earth have been found. But many pieces

of the jigsaw puzzle remain to be put in place.

Advances of great importance in ionospheric

radio communications and in meteorology will

undoubtedly appear when we have a satisfactory theoretical explanation that fits together

the many known solar-terrestrial relationships.

A specific example may help to illustrate

the expected course of progress. We now know

that world weather patterns sometimes respond

in a rather clear-cut way to changes in the

character and level of solar activity as measured,

for instance, by the average intensity of the

emission of the sun's green coronal line. Although our understanding of the solar physics

of the situation is very imperfect, we infer that

the observed large changes in the coronal

emission probably signify large changes in the

total ultraviolet or X-ray emission of the sun.

Nonetheless, the total changes of energy of

the sun's emission must be very small, probably

measured in thousandths or millionths of the

solar constant. Not only do the world weather

changes require vastly greater energy than the

sun's variations can supply but terrestrial

physics suggests that the level at which the

X-ray or ultraviolet energy is absorbed must

be so high that it is difficult to imagine that it

could exert a significant influence all the way

down into the weather sphere, far below. Yet

the problem stands, a challenge and a mystery.

No one can doubt that real progress in forecasting will come when we have a satisfactory

physical explanation of the link, whether it

be a trigger mechanism for amplifying infinitesimal solar changes or some yet unknown cause.

The multitude of phenomena of the ionosphere and the many variations of the effects

with latitude and longitude are equally enigmatic problems. Few astronomers or geophysicists today doubt that the earth is subject to

highly irregular and unpredictable "showers" of

solar corpuscles, or that these corpuscles control

important effects in aurorae, the ionosphere, and

earth magnetism. Yet, in spite of the brilliant

victories of the Chapman-Ferraro theory of such

effects, we know that the theory is only a crude

beginning, and that it is grossly inadequate to

account for major phenomena that are well observed. We particularly need new research to

improve our predictions for communications in

polar latitudes.

Many simple, but major, questions of solarterrestrial relationships are yet unanswered.

For example, do cosmic rays come from the sun ?

A few short years ago it was hard to find a working cosmic-ray physicist who would admit that

the sun as much as gently modulated the cosmic

rays that were believed to come in all directions

from outer spaqe. Astronomers generally assumed, without any serious doubts, that per-
NEW HORIZONS IN ASTRONOMY 101

haps well over half of the energy of the universe

was bound up in these mysterious high-speed particle emissions. Now the view prevails that a

major part of the low-energy cosmic rays come

from, or are at least strongly controlled by, the

sun; and how many of the more energetic cosmic

rays also come from nearby, rather than from

the farthest reaches of space, is an open question.

The observational-theoretical picture of solar

activity needs integration. The vast complex

of sunspots, solar magnetic fields, flares, plages,

prominences, and the corona needs to be welded

into a consistent and understandable unit.

From such a unified picture of solar activity,

and of the quiet sun as well, will come better

understanding of the sun's radiational and corpuscular emissions, particularly those emissions

stopped high in the earth's atmosphere. At

the moment, it is hard to say whether the solar

physicist or the rocket and satellite launchers

will be first to specify the nature of the sun's

variable emissions as seen from an altitude, let

us say, of 250 kms. In any event, the advances in the decades immediately ahead will

be spectacular, and the practical stakes in longand short-range weather forecasting may well

surpass the very important gains that can

reliably be expected in the field of communications forecasting.

Solar Physics

By Leo Goldberg1 and Donald H. Menzel3

At the present time, astronomy is in a remarkable period of rapid growth wherein the

rate of accumulation of knowledge is higher

than at any time in the past. In such a

dynamic situation, existing facilities rapidly

become obsolete and the new instruments

required for continued development of the

science are generally larger and more expensive

than their predecessors. Unfortunately, the

traditional sources of financial support of large

astronomical facilities have been seriously

diminished. Few private donors are in a

position to contribute gifts of the magnitude

necessary for modern research installations,

and the larger foundations appear to be concentrating their primary efforts in other areas.

Thus, apart from occasional gifts or contracts

from industrial concerns, the burden of support

must fall upon Government agencies, which

can function wisely only if they are guided by

the collective, considered judgment of responsible scientists. This means that astronomers as a group must assume the responsibility for deciding which large-scale

facilities are urgently needed for the continued

development of their science, and they must

also solve the organizational problems concerned with the cooperative use of these

facilities.

We do not imply that all future astronomical

research should be directed by planning committees. On the contrary, we strongly believe

that the best new ideas and the initiation of

new fields of investigation will come from

individuals or small groups of investigators

who must continue to receive the strongest

possible financial support from the National

Science Foundation, the Office of Naval Research, and other Government agencies. The

regular program of grants and contracts by

' The Observatory, University of Michigan, Ann Arbor, Mich.

• Harvard College Observatory, Cambridge. Mali.

such agencies supports a wide variety of studies

and is frequently responsible for so-called

"seed" research, in which the expenditure

of relatively few thousands of dollars may lead

to the germination of a whole new field of

study. The next step in the nurturing of the

seedling may well require the construction of

equipment costing hundreds of thousands or

even millions of dollars. It is at this point

that planning on the national level must enter

to evaluate the potentiality of the study with

reference to astronomy as a whole. It would

be unfortunate indeed if the question of cost

should be a major factor in any decision against

embarking in some new, exciting, and otherwise

worthwhile project. This country is capable

of extending far greater financial support than

it has in the past to provide an enormously

expanded effort in basic research. The really

serious bottleneck in astronomy is the current

shortage of topflight astronomers. We must

therefor e guard against wasting precious

scientific manpower in large but relatively

unproductive projects, especially those financed by public funds. The shortage of

astronomers is a national problem which should

be a major concern of the National Science

Foundation.

No branch of astronomy has progressed more

rapidly since 1940 than has solar physics.

We do not except radio astronomy, which, as

a special technique of tremendous power

rather than a separate branch of astronomy,

has supplemented optical methods in contributing importantly to virtually every field of

solar research.

Following a period of phenomenal growth at

the beginning of the century, sparked by the

invention of the spectroheliograph and the solar

tower and by improvements in spectroscopic

gratings, solar research progressed rather slowly

through the 1920's and early 1930's. Largely

103
104 SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS

through the efforts of George Ellery Hale,

observational solar physics had advanced much

more rapidly than the theoretical interpretation

of the data. This gap, however, began to close

in the late twenties, with the development of

theories of ionization equilibrium and of atomic

structure.

Although by 1930 a certain number of successes had been achieved from the application

of the new physical theories to the observations, they were far outnumbered by many

conspicuous failures. In 1930, numerous roadblocks filled almost every avenue of solar

physics. It is useful to survey these barriers

from the vantage point of 1956, both to note

those that have already been removed, or are in

the process of being eliminated, and to assess

how far we may now expect to move down the

road before new obstacles appear.

The photosphere

One of the most fundamental problems in

solar physics concerns the structure and

circulation of the gaseous layers comprising

the photosphere, the region in which the

continuous spectrum is formed and in which

most of the Fraunhofer absorption takes place.

The solution of this problem calls for the construction of a model photosphere in which

the chemical composition and the distribution

of temperatures, pressures, gas velocities, and

magnetic fields are completely specified everywhere in the photosphere. In theory, we can

derive all of this information from observations

of the continuous spectrum and the Fraunhofer

lines. Practically, however, the problem is

frightfully complicated. For example, measurements of solar limb darkening and spectral

energy distribution will yield the distribution

of temperature with depth if we know what

processes produce the continuous spectrum.

We can then calculate the pressure distribution

if the atmosphere is in hydrostatic equilibrium.

This assumption is certainly wrong if the atmospheric gases are in violent, turbulent motion

caused by convective instability or electromagnetic forces. To construct an accurate

model we must evaluate the relative contributions of radiation and convection in the

transport of energy through the atmosphere.

Stability depends on the temperature gradient

and the chemical nature of the atmosphere.

The temperature gradient, in turn, must be

sufficient to drive the solar radiation through

the semiopaque gas. In an idealized atmosphere, where the atmospheric gases are artificially stirred and where no heat is transferred

from or into the moving volume, the temperature gradient must follow the so-called adiabatic

law. Such an atmosphere is in neutral equilibrium. Exchanging equal masses of gas at

different atmospheric levels produces no effect

on the temperature distribution. As the mass

from the lower level rises and expands, it

assumes the same volume, pressure, and

temperature as that of the mass it replaces.

As long as the temperature decreases upward

at a rate slower than that specified by the

adiabatic condition, the atmosphere will be

stable. Radiative transfer will predominate

and the flux of energy will be essentially constant over the entire solar disk. However, if

the natural gradient is steeper than the

adiabatic, a mass of gas displaced upward

becomes warmer than its surroundings. It will

continue to rise, like a hot-air balloon. Analogously, a mass displaced downward becomes

cooler than its surroundings and thus descends.

An atmosphere, so constituted, becomes unstable and convection sets in.

When convection is mild, the energy conveyed upwards will essentially equal that conveyed downwards. Transfer by radiation may

still predominate. However, as convection

increases, nonlinear effects enter so that we no

longer have a precise balance. We shall then

have a net outward transport of energy by

moving airmasses. The convective and radiation fields may be said to "couple" together.

The radiative flux at the outer boundary will

vary from point to point over the solar disk in

such a way as to maintain a constant average.

The convective field will carry momentum as

well as energy and thus will produce either diminution or increase of the effective gravity, with

consequent distention or compression of the

atmospheric gases.

Unsold has shown that convective instability

tends to exist in two separate levels of the solar

atmosphere, one below the photosphere and the

other near the top of the reversing layer. The

lower one arises from ionization effects of hvdro-
NEW HORIZONS IN ASTRONOMY 105

gen, the dominant chemical constituent of the

solar atmosphere. If a volume consisting

largely of ionized gas is displaced upward, it

expands and cools. The cooling effects lead to

electron capture. The released energy raises

the temperature above the value prescribed by

the adiabatic relation, and thus we find a hydrogen convective zone. The second zone is

associated with the region where the atmosphere

is transparent except in the absorption lines.

The "blanketing effect" of the absorbing atoms

again changes the temperature gradient in

such a way as to induce convective instability.

When we turn to the Fraunhofer spectrum for

information on chemical composition and gas

motions, we encounter even more difficult

problems. The shape and intensity of a solar

absorption line depend in a very complicated

manner on a variety of factors, of which the

abundance of the responsible element is only

one. The "f-value," "line strength," or absorptive power of an atom for a given radiation

is also significant. The population of atoms in

the two levels responsible for the line formation

is also extremely important. The Boltzmann

equation fixes these populations if the medium

is in thermodynamic equilibrium. The temperature is the determining parameter. Likewise,

the Saha equation, under similar conditions,

fixes the amount of dissociation. However, the

very fact that energy is flowing through the

medium requires departures from this equilibrium condition. An exact determination of the

significant populations theoretically requires

advance knowledge of all the atomic constants,

such as areas for absorption of radiation or

collisional excitation, and the appropriate

transition probabilities. The equations of statistical equilibrium are most cumbersome and

difficult to solve. Further progress would

seem to depend on the development of new

techniques for the obtaining of approximate

solutions near thermodynamic equilibrium.

We need to know the mechanism that the

atoms employ for removing energy from the line.

The process may involve pure absorption,

coherent scattering, or noncoherent scattering,

and it is usually not obvious which mechanism

or combination of mechanisms operates for a

given line at a given height in the atmosphere.

The character of the line depends critically on

the assumed atmospheric model. Third, we

must know or determine what processes cause

the broadening of the spectral lines. The most

important broadening mechanisms include random Doppler effects caused by thermal agitation or by macro- or micro-turbulence, collisions, Stark effect, and radiation damping.

Widening from hyperfine structure and Zeeman

effect may not always be negligible. The

parameters necessary for calculating the Stark

and collisional broadening are very difficult to

come by, theoretically or experimentally.

Twenty-five years ago only a rough beginning

had been made towards the quantitative interpretation of the solar spectrum. Virtually all

analyses of the Fraunhofer spectrum were

based on the idealized assumption that the

continuous spectrum originated in a sharply

bounded radiating surface, the photosphere, and

that the lines were produced in a cooler superimposed layer, the reversing layer, which was

supposed to be at constant temperature and

pressure. Only very crude guesses were available for most f-values and line-broadening

parameters, and accurate intensity measurements had been made for but a mere handful of

solar lines. Nevertheless, H. N. Russell (1929)

accomplished a remarkable feat in achieving

what he has called the "zero-th approximation"

to the chemical composition of the sun, especially in view of the fact that only recently have

astrophysicists displayed any general agreement

that the "first approximation" has finally been

reached.

No one doubted the existence of large temperature and pressure gradients in the solar

photosphere, but attempts to construct a model

solar atmosphere were frustrated by a total lack

of knowledge concerning the origin of the continuous spectrum. Calculations by Menzel and

Pekeris (1935) gave too low a figure for the

absorption by transitions involving the negative

hydrogen ion. In 1938 Wildt (1939) noted the

existence of a bound state of H —, and suggested

that photoelectric ionizations and free-free

transitions of this negative ion were the major

causes of the sun's continuous opacity. Wildt's

suggestion was confirmed by the work of

Chandrasekhar and associates (Chandrasekhar

and Breen, 1946) and Chalonge and Kourganoff

(1946).
106 SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS VOL. 1

Numerous calculations of model solar photospheres followed. Even now, we must admit

that the photospheric model is known only in

the first approximation. Although most

workers agree on the temperature distribution

in the intermediate layers, the structure of the

deepest and uppermost layers remains uncertain. Since most of the radiation from these

regions is absorbed by the overlying gases, we

must use theory rather than direct observation

for the analysis. In consequence, the location

and extent of the convective zone are not

precisely known. Indeed, the magnitude of

what appear to be convective processes in the

outer atmosphere suggests that even the intermediate layer may not be free from effects of

the hydrogen zone. Radiation from the upper

layers can best be observed at the extreme limb

of the sun, where terrestrial atmospheric seeing

hampers observation. The blanketing effect of

the absorption lines and possible departures

from thermodynamic equilibrium further complicate the problem.

Additional difficulties of even more basic

character stand in the way of a definitive model

of the photosphere. Just a few years ago it

seemed possible that a straightforward analysis

of measurements of solar limb darkening and

absolute spectral energy distribution would determine an accurate solar model. It was hoped

that the model could be tested and further refined by calculation of absorption-line profiles

in terms of a specific chemical composition of

the atmosphere. For several reasons, this hope

now seems to have been optimistic. First, no

existing model has succeeded in reproducing in

detail the profiles of strong absorption lines. In

particular, the observed central intensity has

always been greater than the calculated, but the

discrepancy has usually been attributed to inperfections in the theory of line formation.

Various devices such as fluorescence effects and

noncoherent scattering have been employed to

fill in the line center, but none has proved quite

satisfactory. Second, the empirical dependence

of the continuous opacity upon wavelength, derived from Pierce's (in press) latest measures of

limb darkening, differs significantly from the

values calculated in terms of currently accepted

electron pressure distributions and the theory of

continuous absorption by neutral hydrogen atoms and negative hydrogen ions. Errors evidently exist either in the model or in the theory

of the continuous absorption, or perhaps in both.

Third, no spherically symmetrical model of the

photosphere appears capable of representing the

center-to-limb variation in the intensities of

both low and high excitation lines of Fe I

(Bohm, 1954). Nor can the center-limb variation in the CO line intensities be accounted for

on such a model (Newkirk, 1953). Similar results have been obtained by Elste (1955) at

Gottingen and by others.

The evidence is accumulating that the lineintensity observations can be reconciled with

theory only when account is taken of temperature and velocity fluctuations in the photosphere and chromosphere. The recognition of

temperature fluctuations is not new, since they

are reflected in the brightness fluctuations of

the granules, but their influence on the Fraunhofer spectrum represents a second-order effect

which could not be detected before the advent

of modern techniques of observation and analysis. It is now generally agreed that the phenomenon of the granules is a consequence of

turbulent convection in the photosphere and

that, on the average, the bright and presumably

hot granules are ascending and the cooler, intergranular regions are descending in the photosphere. If the convection is sufficiently violent,

however, coupling of the mechanical and radiative fields may occur, with nonlinear effects.

The optical thickness of the atmosphere might

vary significantly between the bright and dark

areas, with an appreciable change of line intensity.

Another consequence of the turbulence is the

excessive broadening of the centers of the

Fraunhofer lines as compared with that to be

expected from thermal Doppler effect at a

temperature of 6,000° K. The root-meansquare turbulent velocity derived from individual line profiles and from the curve of growth

is about 2 km/sec. Richardson and Schwarzschild (1950) supposed that the radial velocities

of individual granules could be obtained from

the measurement of minute Doppler shifts from

point to point along the length of Fraunhofer

lines, under conditions of high spectroscopic

resolution and exceptional seeing. Careful

measurement of selected Mount Wilson plates
NEW HORIZONS IN A8TRONOMY 107

did indeed reveal irregularities in the lines,

which, upon measurement, yielded a root-meansquare velocity of about 0.4 km/sec. This effect

was interpreted as arising from the macroturbulence of relatively large elements, the higher

velocity of 2 km/sec being attributed to the

microturbulence of elements smaller than the

thickness of the "reversing layer."

It had been supposed that the observation

of the random Doppler shifts was limited by

the resolution of the solar image; i. e., by the

external seeing. Recent work by McMath

(1956) and McMath, Mohler, and Pierce

(1955) at the McMath-Hulbert Observatory

has revealed, however, that limitations in

spectroscopic resolution are at least of equal

importance and that air turbulence within the

spectrograph may "wash out" the small

Doppler effects. On McMath-Hulbert photographs made with the vacuum spectrograph,

all photospheric lines show a characteristic

"zigzag" appearance even when the seeing

quality is less than average. Further, when

the seeing is good, the lines fluctuate in intensity from point to point along the solar

disk. Thus the temperature and velocity

fluctuations in the solar atmosphere may now

be directly observed in the Fraunhofer spectrum,

at least for elements about 2 seconds of arc in

diameter or larger.

Several inferences may be drawn from these

results, the most general of which is that the

interpretation of photospheric and chromospheric observations on the basis of spherical

symmetry should now be abandoned, and the

observational emphasis directed towards determining the nature of the velocity and

temperature fluctuations in the solar atmosphere. The impact of this new development upon the National Science Foundation is

likely to be very great. First, important funds

will be required both to modernize existing

solar spectrographs and to construct new ones.

Second, an enormous quantity of new and

valuable data will become available for analysis,

which will necessitate funds for assistance in

reduction and analysis. Third, new impetus

will be provided for the development and

refinement of image-tube techniques. At present photographic exposure times are not less

than 10-20 seconds, whereas the image tube

may be expected to reduce exposures to a small

fraction of a second. Such a reduction would

greatly reduce the effects of atmospheric

scintillation and improve our determinations

of the lifetimes of small-scale solar phenomena.

Finally, it is already apparent that spectroscopic

resolution has far exceeded the image resolution

and that limitations in external seeing quality

will seriously hinder the discovery of the details

of solar atmospheric structure.

Poor external seeing is caused both by disturbances inherent in the atmosphere and by

local disturbances of the air produced by solar

heating of the telescope and mirrors. The National Science Foundation panel for a national astronomical observatory, under the chairmanship

of R. R. McMath, has recently considered this

problem and has concluded that two courses

of action should be taken. First, the present

site survey for the national observatory should

be extended to include observations of daytime

as well as nighttime seeing. Second, an investigation of instrumental seeing should be

undertaken at the earliest possible moment

with the goal of designing a solar telescope in

which heating effects would be minimized.

It is probable that the solution of the problem

of telescopic seeing is now the most important

instrumental challenge in the field of solar

physics.

The vacuum spectrograph or other spectrographs that may be designed to eliminate

instrumental turbulence will initiate a new

phase of research in solar spectroscopy that will

take many years to run its course. The end

result should be a much finer elucidation of

the structure of the solar atmosphere, including

its composition and circulation, than is possible

at present. As a valuable byproduct we may

also expect a better understanding of the

mechanisms of absorption-line formation. As

previously mentioned, the fact that the observed

central intensity of a very strong line is always

greater than that calculated from theory has

been a major difficulty. However, it is already

obvious from the vacuum spectrograms that

some if not all of the discrepancy between

theory and observation will vanish when we

allow for fluctuations in observed central

intensity, which have generally been blurred

by poor seeing inside the spectrograph. In
108 SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS VOL. 1

addition, theorists will have to include effects

of mechanical as well as radiations! transfer

of energy.

Finally, it may not be out of place to point

out that if temperature fluctuations are present

in the solar photosphere they undoubtedly occur

also in stellar atmospheres, especially in the

highly turbulent atmospheres of the supergiant

stars.

Chromosphere and corona

By definition, the beginning of the chromosphere is the upper boundary of the photosphere.

It is also the transition region, perhaps 20,000

km thick, in which the electron temperature

increases from its photospheric boundary value

of 4,000°-4,500° K to 1,000,000° K or higher, in

the inner corona. It is therefore of strategic

importance, since within its confines probably

occur the processes that maintain the corona at

its high temperature. More than 20 years ago

Menzel (1931) called attention to the anomaly

posed by the coexistence in the flash spectrum

of emission lines of both neutral metals and

ionized helium. Menzel pointed out that the

observed intensity of He+X4686 could not be

explained by a temperature of 6,000° K even if

the chromosphere were composed of pure helium, and suggested that the required ionizing

radiation might come from a hypothetical

ultraviolet excess in the solar spectrum.

The problem posed by Menzel is still paramount for theories of chromospheric structure,

even though the solution in terms of an ultraviolet excess has been ruled out. Since the

1930's, at least three important developments

have paved the way for a renewed attack on the

chromospheric problem. First, the discovery

of the high-temperature corona inevitably requires the existence of large temperature gradients somewhere in the chromosphere. Second,

the observation from rockets of the far ultraviolet solar spectrum has shown that the hypothetical ultraviolet excess does not exist.

Third, the observations clearly show that the

chromosphere is not a static phenomenon and

that hydrodynamic or aerodynamic processes

play an important role in fixing its physical

state.

Evidence bearing on the structure of the

chromosphere comes from observation of the

optical spectrum during eclipses, from motion

pictures in Ha obtained with large coronagraphs, from observation of the radio emission

at high frequencies, from observation of the

central intensities of certain strong Fraunhofer

lines and of the XI0830 absorption line of He I.

from observation of emission lines in the rocket

ultraviolet, and from observations of the earth's

ionosphere. Interpretation of the observations

from the conventional point of view of spherically symmetrical chromospheric models leads

to serious discrepancies. In particular, one set

of observations has seemed to require a temperature for

 the low chromosphere in the neighborhood of 5,000°-6,000° K, whereas another set

seems to demand much higher temperatures of

about 20,000° K or greater.

A possible path out of the dilemma was

pointed out in 1949 by Giovanelli (1949), who

suggested that the chromosphere is nonuniform

and that it consists of alternate "hot" and

"cold" columns. Hagen's (1954) measurements

of the center-limb variation of the radio emission also point to a model of this general type.

Strong support for the nonuniform model has

recently been provided by the detailed analysis

of flash spectra by the Harvard-High Altitude

Observatory group (Athay, Evans, and Roberts,

1953; Athay, Billings, Evans, and Roberts,

1954; Athay and Roberts, 1955; Pecker and

Athay, 1955; Matsushima, 1955; Athay, Menzel, Pecker, and Thomas, 1955; Athay and

Thomas, 1955, 195Ga, 1956b; Athay and Menzel, 1956; Billings, Cooper, Evans, and Lee,

1954) and by studies of the structures of

Fraunhofer lines at the McMath-Hulbcrt Observatory (McMath, Mohler, Pierce, and Goldberg, 1956; Mohler and Goldberg, 1956).

Although the exact geometry and physical

properties of the hot and cold regions have

not yet been clearly established, interesting

working hypotheses have been suggested whicli

clearly indicate the lines along which future

researches should be planned. Observations of

the flash spectrum at future eclipses will continue to play a leading role in the investigations

of the chromospheric structure, but they may

now be supplemented by additional new techniques which have recently been developed.

Of great importance are observations of microwave radio emission, which should be carried
NEW HORIZONS IN ASTRONOMY 109

out with large antennae of the highest possible

resolving power. The accurate measurement

of emission profiles of solar lines in the rocket

ultraviolet will also be crucial for the testing of

chromospheric models. Finally, the Fraunhofer spectrum itself constitutes a rich but

largely neglected source of chromospheric

information. The interpretation of the limb

spectrum is complicated by the two-dimensional

integration, both radial and tangential, of the

chromospheric emission. On the disk, the line

of sight integration is only one dimensional.

Indeed, the complex structures of the X10830

helium line and of the cores of the hydrogen and

ionized calcium lines, as revealed by the vacuum

spectrograph, are undoubtedly a consequence

of the temperature fluctuations in the chromosphere and should be analyzed from this point

of view.

The dynamic nature of the chromosphere is

graphically portrayed by the Ha motion pictures recently obtained by R. B. Dunn of

Harvard and the Upper Air Research Observatory with the 15-inch Sacramento Peak coronagraph. These remarkable photographs record

the rapid changes that occur in the irregular

structure of the outer chromosphere. Since

the chromosphere is basically a dynamic phenomenon, with kinetic transport of both

energy and momentum, a true model of this

region of the solar atmosphere should depict

and account for the internal motions. One

may question whether any static solution of

the fundamental equations exists. Small perturbations about such a hypothetical static

solution, for example in terms of a linear

combination of normal modes, seems out of

keeping with the actual observations. Instead,

we suspect that the mean energy transport,

determined by the solar energy generation,

can never be fulfilled in a purely radiative

atmosphere.

The solution will depend on the magnetohjdrodynamic equations of motion in their

nonlinearized form. The Bhatnagar-GrossKrook (1954) equation appears to be a useful

substitute for the much more complex Boltzmann transport equation. The simplification

represents a new approach to problems of gas

dynamics in which the distribution of velocities

no longer follows Maxwell's law. Instead,

the equation represents the velocities as lying

between the two extremes of isotropic scattering

without change of velocity, or of reemission

under thermodynamic equilibrium at the temperature described by the root-mean-square

velocity. The equation bears a close resemblance to the Milne-Eddington equation of

radiative transfer, with solution between the

extremes of monochromatic radiative equilibrium and local thermodynamic equilibrium.

The solution of the transfer problem, especially when electric currents and magnetic

fields are present in the medium, can rarely

be expressed in terms of known functions.

Instead we shall have to determine families of

solutions, with ranges of boundary values and

initial conditions. In general, these calculations will require the use of high-speed computing machines. The National Science Foundation may well be called on to furnish funds for

the development of models of the dynamic

chromosphere.

Closely associated with structural problems

of the solar atmosphere is the question of origin

of the solar radio emission. It is well known

that an ionized gas, with specified electron

density, possesses the so-called "plasma" oscillation frequency equal to the "critical frequency"

at which the phase-velocity becomes equal to

zero. Many persons have supposed that natural oscillations of such a medium will lead to

the emission of radio waves on the plasma

frequency. A detailed analysis indicates, however, that the plasma oscillations, which resemble sound waves in their longitudinal

character, cannot be converted into transverse

waves by means of boundary conditions alone.

In short, plasma oscillations cannot produce

transverse radio waves. When the medium

contains a magnetic field, however, the rigidity

of the medium associated with magnetism

resolves any longitudinal wave into vibrations

having one transverse component. A disturbance propagated from regions of high into

regions of low density tends to have its vibrations amplified. In this way, as Krook and

Menzel have shown, a magnetohydrodynamic

wave at low density will eventually escape from

the sun's outer boundary as a true electro-
110 SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS V«L.

magnetic wave, with polarization determined

by the orientation of the internal magnetic

field.

Low-frequency radiation tends to be trapped

or attenuated in an atmosphere containing

ions and electrons. Thus, only higher frequencies are able to penetrate from the chromosphere through the overlying coronal regions.

Similarly, a given disturbance rising through the

solar corona permits lower and lower radio

frequencies to escape. Study of the changing

spectrum of solar radio noise, especially in

combination with optical data, should give

valuable information about the solar atmosphere. J. P. Wild's (1953) dynamic spectrum

analyzer presents the observational facts in

graphic form. The Cornell measures of solar

radio noise, as analyzed by H. W. Dodson

(1953), have already yielded significant correlations with flare and other solar activity.

The corona itself poses many difficult problems. There is little doubt that the streamers,

domes, and arches owe their basic structure to

general and local magnetic fields. Magnetic

focusing must be important in certain of these

phenomena. The forces associated with electric

currents undoubtedly play a significant role in

determining the motions of prominences and

the general condensations that occur from the

solar corona. The detailed theory of fluid gas

dynamics necessary to account for these phenomena will undoubtedly be very complicated.

Solar activity

The term "solar activity" includes all types of

phenomena of a transient nature except the

small-scale fluctuations in velocity and temperature associated with granulation. Typical of

these phenomena are sunspots, faculae, plage

regions, flares, filaments, prominences, and

coronal disturbances. Solar activity occurs

throughout the solar atmosphere; e. g., the sunspots and faculae in the photosphere, the flares

in the chromosphere and corona, and the prominences largely in the corona. No attempt will

be made in this brief survey to summarize the

problems related to solar activity in any systematic way. The entire subject is still in a

highly chaotic state, hardly less so at present

than in the 1930's. One reason for this state

of affairs is the aforementioned inherent diversity and complexity of solar phenomena. A

second reason is that only in rare cases have the

most modern spectroscopic techniques been applied to the study of solar details. The changes

on the solar surface are so gross and easy to

observe that solar observers have been tempted

to make large numbers of observations with instruments of relatively low resolving power,

either visually or on uncalibrated plates.

The rapid acceleration of observing programs

during the International Geophysical Year gives

to those participating in the program the responsibility of maintaining high standards for

their records. It is to be hoped that the need

for large masses of data will not distract the

attention of solar physicists from the really basic

problems of solar activity, which pertain to the

origin and physical nature of solar phenomena.

In the end, a complete understanding of the

effects of solar radiation on the earth is more

likely to come from the solution of these basic

problems than from correlations developed between one or two observed solar quantities and

various geophysical data.

Knowledge of the physical state of the various

kinds of solar phenomena must precede correct

theories of their origin. In the light of recent

technical developments—which include the perfection of gratings and the magnetograph by

the Babcocks, the construction of the vacuum

spectograph by McMath, and the development

of a large coronagraph by the Harvard-Air

Force group on Sacramento Peak—we can look

forward to a fruitful period of quantitative

analysis of the spectra of solar details during the

forthcoming sunspot maximum.

An understanding of the detailed structure of

sunspots, including the three-dimensional distribution of temperatures, pressures, gas velocities, and magnetic fields, should be a primary target during this period. There also

seems to be no reason why high-dispersion

studies of flare spectra should not lead to know1-

edge of the vertical temperature and pressure

distributions in the flare-emitting layers, which

can provide insight into the mechanisms of

origin both of the flares themselves and of the

ultraviolet radiation responsible for ionospheric

disturbances. In this connection we consider

that the physical investigation of faculae and

the associated "plages faculaire" has been
NEW HORIZONS IN ASTRONOMY 111

grossly neglected. It is well established that

flares break out only in regions occupied by the

plages. The latter, which are also regions of

abnormal magnetic activity, are generally located above the photospheric faculae. There

must exist, then, certain physical properties

associated with faculae and plages that favor

the outbreak of flares; yet the properties of

these regions have hardly been studied.

Our understanding of the physics of prominences is equally weak. The development of

the motion-picture technique by McMath and

by Lyot in the early 1930's has been followed

by large numbers of reasonably accurate measurements of prominence motions, chiefly in Ha

radiation. Some of the observed motions have

been explained plausibly, but nevertheless only

qualitatively, on the basis of the combined

laws of hydrodynamics and electrodynamics by

Alfv6n and by Menzel, but a detailed quantitative description of prominence behavior is still

somewhere in the future. Here again, the

development of a successful theory is retarded

by uncertain knowledge of the temperatures,

pressures, turbulent motions, and magnetic fields

within the prominences. The first three of

these quantities are readily within the reach of

modern instruments and theory. But one must

avoid the temptation to tackle these important

problems with equipment yielding marginal

data that can be interpreted only on the basis

of doubtful assumptions or even wishful

thinking.

Finally, although we have devoted major

attention to observational trends in solar

physics, we have not wished to minimize the

importance of support by the National Science

Foundation and other Government agencies for

theoretical research of varying shades of purity.

The so-called theoretical projects that are most

likely to receive support today are those that

involve detailed computations— usually by highspeed calculators—based on theories already in

existence. In short, practical results are guaranteed, but such projects are no more theoretical

than those that involve the measurement of

stellar magnitudes or of spectral line intensities.

We do not question the value of accurate computations to astronomy and solar physics.

However, the National Science Foundation

should also consider how its support might

stimulate the birth of new theories as well as

foster applications of existing ones.

For example, the solution of many problems

of solar structure or solar activity waits upon

further developments in the field of generalized

gas dynamics concerning the motions of hightemperature gases in the presence of various

types of force fields. Indeed, the United States

appears to be lagging in this important field,

which has applications to physics, geophysics,

and industry as well as to astrophysics.

Several symposia on this general subject have

already been held and have been useful both in

terms of summarizing what little knowledge

exists in this area and in stimulating new

thinking. If a really top-notch group of investigators could be assembled to work together

for a relatively long period—for months, or even

years, instead of a few days—to study the

problem in its broadest aspects, we could

expect significant new advances. For a problem of this magnitude, support would be

necessary for a number of years, during which

headquarters of the project would be maintained at a particular university with some

rotation of personnel. A center of this sort

devoted to theoretical problems of fluid gas

dynamics could be as important to astronomy

and to other branches of science as a cooperative

observing facility.

References

ATHAY, R. G., BILLING*, D. E., EVANS, J. W., AND

ROBERTS, W. O.

1954. Astrophy.s. Journ., vol. 120, p. 94.

ATHAT, R. G., EVANS, J. W., AND ROBERTS, W. O.

1953. Observatory, vol. 73, p. 244.

ATHAT, R. G., AND MENZEL, D. H.

1956. Astrophys. Journ., vol. 123, p. 2S5.

ATHAT, R. G., MENZEL, D. H., PECKER. J.-C, AND

THOMAS, R. N.

1955. Astrophys. Journ. Suppl., vol. 1, p. 505.

ATHAT, R. G., AND ROBERTS, W. O.

1955. Astrophys. Journ., vol. 121, p. 231.

ATHAT, R. G., AND THOMAS, R. N.

1955. Astrophys. Journ. Suppl., vol. 1, p. 491.

1956a. Astrophys. Journ., vol. 123, p. 299.

1956b. Astrophys. Journ., vol. 123, p. 309.

BHATNAQAB, P., GROSS, J., AND KROOK, M.

1954. Phys. Rev., vol. 94, p. 511.

BILLINGS. D. E., COOPER, R. H., EVANS, J. W., AND

LEE, R. H.

1954. Electronics, vol. 27, p. 174.

B6HM, K.-H.

1954. Zeitschr. Astrophys., vol. 35, p. 179.
112 SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS

CHALONGE, D., AND KOURGANOFF, V.

1946. Ann. d'Astrophys., vol. 9, p. 69.

CHANDRASEKHAR, S., AND BREEN, F.

1946. Astrophys. Journ., vol. 104, p. 430.

DODSON, H. W.

1953. In Kuiper, ed., The sun, p. 692.

ELSTE, G.

1955. Zeitschr. Astrophys., vol. 37, p. 184.

GIOVANELLI, R. G.

1949. Monthly Notices Roy. Astron. Soc., vol.

109, p. 293.

HAGEN, J. P.

1954. Jouru. Geophys. Res., vol. 59, p. 158.

MATSU8HIMA, S.

1955. Astrophys. Journ. Suppl., vol. 1, p. 479.

MCMATH, R. R.

1956. Astrophys. Journ., vol. 123, p. 1.

MCMATH, R. R., MOHLEB, O., AND PIERCE, A. K.

1955. Astrophys. Journ., vol. 122, p. 565.

MCMATH, R. R., MOHLER, O., PIERCE, A. K., AND

GOLDBERG, L.

1956. Astrophys. Journ., vol. 124, p. 1.

MENZEL, D. H.

1931. Publ. Lick Obs., vol. 17.

MENZEL, D. H., AND PEKERIS, C. L.

1935. Monthly Notices Roy. Astron. Soc, vol.

96, p. 77.

MOHLER, O., AND GOLDBERG, L.

1956. Astrophys. Journ., vol. 124, p. 13.

NEWKIRK, G.

1953. Thesis, Univ. Michigan.

PECKER, J.-C, AND ATHAT, R. G.

1955. Astrophys. Journ., vol. 121, p. 391.

PIERCE, A. K.

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RICHARDSON, R. S., AND SCHWARZSCHILD, M.

1950. Astrophys. Journ., vol. Ill, p. 351.

RUSSELL, H. N.

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WILD, J. P.

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WILDT, R.

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Stars and the Galaxy

Spectral Classification

By W. W. Morgan *

In a classical review titled "Some Problems

of Sidereal Astronomy" prepared by H. N.

Russell in 1917, the main object of astronomy

is stated as being

. . . the development, on the basis of collected facts,

of satisfactory theories regarding the nature, mutual

relations, and probable history and evolution of the

objects of study.

Russell suggests that before the stage of

formation of a general theory is reached, two

courses may be followed:

(1) to collect masses of information, as accurate

and extensive as possible, by well tested routine

methods, and leave it to the insight of some fortunate

and future investigator to derive from the accumulated

facts the information which they contain regarding

the general problems of the science; (2) to keep these

greater problems continually in mind, and to plan the

program of observations in such a way as to secure as

soon as practicable data which bear directly upon

definite phases of these problems.

The present survey gives an approximate

picture of the research situation in spectral

classification in 1956, a comparison with

Russell's survey, and some suggestions of the

probable directions for future work.

The Henry Draper Catalogue, which was

published between 1918-1924, remains the most

important fundamental body of data in Russell's first category. Russell's conclusions are,

to a large extent, based on the Harvard work.

The interval 1917-1955

Two fundamental conclusions were drawn by

Russell from the study of the lines of stellar

spectra: (a) that almost all of the observed

lines are identifiable with known elements;

and (b) that the vast majority of stellar spectra

fall into a single, continuous linear sequence.

Later work has confirmed fully Russell's first

conclusion; the second requires some qualification.

1

 University of Chicago, Chicago, III.

As early as 1897, Miss Antonia Maury had

reached the conclusion that—

. . . a single series was inadequate to represent the

peculiarities which presented themselves in certain

cases, and that it would be more satisfactory to assume

the existence of collateral series.

She devised a system of 22 "groups" which

formed a sequence; most of the "groups" were

then divided into "divisions," which, in turn,

introduced a second dimension. It was from

this pioneering work of Miss Maury that the

field of spectroscopic parallaxes derives, through

the later work of Hertzsprung, Russell, Adams,

and Kohlschutter.

The situation in 1920 was, therefore, approximately as follows: spectral classification was

considered to be primarily a one-dimensional

problem; in addition, differences in luminosity

were observable by means of certain sensitive

line-ratios. The early work of Miss Maury

had not been adopted as standard, nor was it

continued on any appreciable scale.

If a single development could be considered

of primary importance in spectral classification

since 1920, it would be the realization of the

fundamental two-dimensional nature of the

problem. The critical work here was that of

Bertil Lindblad, who developed criteria for

the separation of stars of differing luminosity

by use of plates of low dispersion; this last

property permitted the extension of the spectroscopic parallax method to great numbers of

stars too faint to be observed with higher

dispersion.

Another fundamental development was the

formalizing of the earlier work (in particular,

that of Lindblad) in the two-dimensional

Yerkes system of Morgan and Eeenan. As far

as the actual criteria of classification are concerned, there was little new in the latters'

system; the principal development lay in the

recognition that spectral classification per se

115
116 SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS

is a two-dimensional process and that the

designation of the "luminosity class" should

be by means of an empirical notation similar

in nature to the ordinary "temperature type."

The process of determining spectroscopic

absolute magnitudes was thus divided into two

stages: (1) the determination of a spectral

type and luminosity class; and (2) the calibration of the empirical luminosity classes in terms

of absolute magnitude. The first of these

operations is purely a description of the spectra.

This description is a fundamental property of

the latter, and there are no compelling reasons

why it would have to be revised with the passing

of time. In this respect the second operation

can be improved with the passing of time, due

to more and more accurate calibration data

becoming available.

Another development in the intervening

years has been the devising of methods for

segregating certain kinds of stars by means of

spectra of exceedingly low dispersion. The

most important applications of these methods

have been to the use of objective-prism

spectrograms.

In 1950 attention was called to the possibility

of classifying members of certain "natural

groups" on objective prism spectra. For

example, a single criterion is sufficient for the

segregation of a well-defined group of early A

main-sequence stars; this is carried out by

means of the Balmer series of hydrogen. If a

group of stars is selected having the characteristic of maximum intensity for the Balmer

lines, it is found that the stars in this group

possess the characteristics of closely similar

absolute magnitude and spectral type. Among

the most interesting of the groups identifiable

on objective-prism plates is that of the blue

giants. Nassau and collaborators at Cleveland

have prepared such a catalog of northern blue

giants, and this work has been extended to

southern declinations and fainter northern

stars by Haro and associates at the National

Astrophysical Observatory in Mexico.

Of comparable importance has been the

development of a method of classification of

late M giants in the infrared spectral region by

Nassau and collaborators. A catalog of the

brightest of these objects has been published

recently, but a beginning has scarcely been

made on this most promising field for future

galactic studies.

A most important development has been in

the field of classification of the faint inhabitants

of dust clouds—the T Tauri stars—by Joy

(Mount Wilson), Herbig (Lick), and Haro

(Tonantzintla). The work on these strange

systems of stars is also in its infancy.

Present and future problems

This topic will be considered under three

general heads: (1) astrophysical interpretation;

(2) spectral classification as applied to problems

of galactic structure and stellar evolution; and

(3) instrumental needs.

(1) Astrophysical interpretation.—Very promising beginnings have been made in the past few

years in devising semiquantitative methods of

classification for stars in certain restricted areas

of the HR diagram. Some of these have made

use of measures of line depths and equivalent

widths to furnish more objective data for classification. Of especial importance is the work

of Greaves, Baker, and Wilson at Edinburgh, of

Chalonge and associates in Paris, and of Hossack at Toronto. Accidental accuracy of a high

order has been achieved. The principal difficulty here is probably in devising a frame of reference sufficiently precise to make full use of the

extremely accurate individual results obtained

by these investigators.

It is from investigations such as these that we

will be able to decide how uniquely the stars can

be classified in a two-dimensional array, and

what the relationship is between the "normal"

and the "peculiar" stars. The practical limit

attainable in the accuracy of spectroscopic parallaxes will also be determined from high-dispersion investigations.

(2) Spectral classification as applied to problems of galactic structure and stellar evolution.—

The application of the spectroscopic parallax

method to large numbers of fainter stars is now

quite feasible. The onty thing needed for a

great increase in our knowledge of the arrangement of the stars in space in the neighborhood

of the sun is more investigators; the method and

equipment are now quite adequate for the work.

New programs will probably be of two kinds:

(1) work on the O-associations, galactic clusters,

and other objects associated with spiral struc-
NEW HORIZONS ENT ASTRONOMY 117

ture; and (2) investigations of the amorphous

disk and halo population.

An interesting group for further investigation

in the second category is that of the late M

giants, which are most important for the determination of the kinematical properties of the

disk population. However, for their observation at great distances it will be necessary to

devise infrared techniques for measurement of

radial velocity and absolute magnitude.

Some of the most dramatic developments of

the 30-odd years since Russell's survey have

been in the field of stellar evolution. The

theoretical work in this field, up to 1951, has

been well summarized by Stromgren (1952).

This work, together with later developments,

now makes possible the application of Russell's

second criterion quoted at the beginning of this

summary. It is now possible to formulate

precise observational experiments concerned

with various aspects of the great general

problem of stellar evolution.

One of the most inviting practical problems

of the future is the invention of a system of

spectral classification according to stellar age,

or relative evolutionary state.

(3) Instrumental needs.—The principal instrumental need for present and projected

observational problems in this field is probably

not for new or larger telescopes but for more

efficient spectrographic equipment for use with

existing instruments.

The requirements for such "maximum efficiency" spectrographs have been outlined

clearly by Bo wen (1947, 1952), who has also

designed and supervised the construction of

instruments of this type which are now in use

with the Mount Wilson-Palomar reflectors.

Special attention should be called to the new

60-inch Cassegrain spectrograph designed by

Bowen which has recently been put into operation. The faint limiting magnitude obtainable

with this instrument illustrates how great a gain

in efficiency is possible with reflectors in the

second category of size. Special mention should

also be made of the exceedingly efficient nebular

spectrograph designed and constructed under

the direction of Mayall for the 36-inch Crossley

reflector of the Lick Observatory.

There are now several large and most useful

objective prism Schmidt cameras in operation.

Especially important new instruments of this

type have been put into operation recently at

Hamburg-Bergedorf, the University of Michigan, and Vanderbilt University. Exceedingly

efficient slitless spectrographs have been designed and constructed recently by Herbig

(Lick) and Meinel (University of Chicago).

References

BOWEN, I. S.

1947. Publ. Astron. Soc. Pacific, vol. 59, p. 253.

1952. Astrophya. Journ., vol. 116, p. 1.

STROMGREN, B.

1952. Astron. Journ., vol. 57, p. 65.

381014—56

Stellar Atmospheres

By Marshal H. Wrubel'

The outer region of a star in which the optical

spectrum is formed is referred to as the atmosphere. This definition delimits a shell lying

between the deeper regions, which cannot be

directly observed, and the higher regions, which

contribute very little energy to the optical

spectrum (but which may be all-important in

the radio spectrum). Although it is convenient

to treat this shell separately, it is certainly not

independent of the layers lying below and

above it.

Any discussion of stellar atmospheres logically

begins with the sun. Here is a star on our

doorstep; and the variety and accuracy of the

solar data we can accumulate far surpasses that

for any other star. We may examine the

radiation received from every point on the

sun's disk and for every wavelength that

penetrates the earth's atmosphere.

In spite of centuries of solar observation,

new techniques continue to introduce new

phenomena. To mention only two of the most

interesting: the Babcocks' magnetograph,

which gives a tracing of the atmosphere's

magnetic field point by point; and the McMathHulbert vacuum spectrograph, which has revealed a remarkable complexity of line structure. Both of these innovations are triumphs

of observational technique and will create

challenges to the theoretician.

The theoretical models that have been explored until recently may be characterized as

being in "hydrostatic and radiative equilibrium"; that is, the atmosphere is supported by

hydrostatic pressure and radiation is the sole

means of energy transport. The logical chain

is this: the temperature distribution is found

either theoretically or by limb darkening

measurements, and it depends on the broad

continuous absorption characteristics of the gas.

A complete physical description of the atmos1

 Indiana University, Bloomlngton, Ind.

phere follows from the temperature distribution

by using the hydrostatic equation, the absorption coefficient, the ionization equation, etc.

The resulting "model atmosphere" combined

with the theory of the formation of absorption

lines is used to predict the observed spectrum.

The only adjustable parameters of such a solar

model are the abundances of the elements which

are determined by a comparison with observations.

Although the reasoning is straightforward,

many problems arise even in the static model.

To list only a few:

(1) The temperature distribution of the

upper layers is uncertain because of the

difficulty of observation close to the limb of the

sun.

(2) The abundance of the second most

abundant element, helium, cannot be determined directly since it is effectively "dead

weight" in a low-temperature atmosphere.

(3) The fundamental data of transition

probabilities for many lines are still missing.

The great emphasis on nuclear and molecular

physics in recent years has attracted physicists

away from research in atomic spectra. A

possible solution is for some young astronomers

to learn the laboratory techniques and to do the

work themselves.

(4) The theories of the shape of the lineabsorption coefficient, i. e., broadening, absorption versus scattering and noncoherent processes, are not in satisfactory form. Of particular interest in this connection is the recent use

of a shock tube at the University of Michigan

to obtain laboratory data on broadening.

(5) The problem of predicting the profile

of an absorption line formed when radiation

passes through many layers of different physical

characteristics has been solved only for very

special cases.

(6) The cumulative effect of many absorp119
120 SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS

tion lines (blanketing effect) is difficult to

evaluate.

However, overshadowing all these difficulties

today is the problem of describing a dynamic

atmosphere. The theoretician no longer can

avoid including large-scale motions in his solar

models but the kinematical description is far

from clear. One important source of information is the "granulation"; but the details are

obscured by "seeing" in the earth's atmosphere. It is to be hoped that Schwarzschild's

unorthodox plan to photograph the sun from

above the "seeing" will be successfully carried

out in the near future.

Although simple models (such as those used

by de Jager) involving temperature inhomogeneities in adjacent columns of the atmosphere

may prove serviceable for a while, eventually

we must have a model obeying the hydrodynamical laws (better still the hydromagnetic

laws) and including the transport of energy by

a combination of convection and radiation.

Increasing use will undoubtedly be made of

high-speed electronic computers to obtain

numerical solutions to these difficult problems.

Computers can certainly not create ideas, but

they can explore the consequences of various

assumptions to an extent previously impossible.

It is a challenge to find efficient machine procedures for calculating model atmospheres and

line profiles under a variety of physical conditions. Computer techniques will probably

be second nature to astrophysicists a few

academic generations hence; at the moment,

however, theoreticians have been very conservative in adopting these new methods.

As if the existing observations do not provide

sufficient problems, astronomers have watched

with great interest as rocket spectrographs and

radio telescopes have broadened the limits

of the observable spectrum. Lyman a has at

last been photographed and even soft X-rays

have been detected. Observations such as

these will have a profound effect on future

models of the sun.

There is very little generally accepted theory

concerning the spectacular transient solar

phenomena. Many solar observatories are involved in the accumulation and classification

of data from which theories will eventually

arise, and a highly coordinated effort in this

direction is being planned for the forthcoming

International Geophysical Year. Many problems in this particular field are treated individually in other articles of this volume.

The techniques developed to describe the

solar atmosphere have been applied with

some success to "normal" atmospheres of other

spectral types. In some cases it is possible

to obtain information that the sun will not

provide directly: for example, the abundance

of hydrogen relative to helium can be found

in principle from the spectra of hot stars.

Unfortunately, no conclusive value can as yet

be given.

Problems immediately arise because of observational limitations. There is just not

enough light to work with. Photoelectric

techniques for the absolute spectrophotometry

of stars are currently being developed by several

observers and should yield important astrophysical data. Nevertheless we have only

indirect measurements of limb darkening, and

these for only a few eclipsing binaries.

As the available light goes down, so does the

possible dispersion, and observations become

less reliable. Some objects are out of reach of

all except the largest reflectors.

Surveys can be carried through by using

rough curve-of-growth methods based on idealized models, but eventually we need detailed

model atmospheres for a wide variety of spectral

types. These are constructed by using theoretical temperature distributions and are tested

by requiring the flux to be constant at every

depth. However, it is difficult to apply this

criterion in early-type stars where the flux

is jagged due to the absorption edges of hydrogen and helium.

As we go toward the cool stars, the importance of the blanketing effect grows and a

premise upon which other models are founded

breaks down—it is not possible to separate

the lines from the continuum. Entirely new

theoretical techniques will have to be combined

with laboratory data on the absorption of molecules before we will have as detailed an interpretation of the atmospheres of stars of late

spectral type as we have for the sun.

Nature has provided a unique tool for studying the outer regions of late-type giants.

Binary systems like 31 Cygni, in which a small
NEW HORIZONS IN ASTRONOMY 121

hot component is gradually eclipsed by the

extended atmosphere of its giant companion,

provide a wealth of detailed information. The

observations indicate a "network of prominences" far more extensive than that of the

sun; and if an understanding of a dynamic

atmosphere is important in the sun, it is vital

here.

As the spectroscopic peculiarities become

more unusual, the extrapolation of normal techniques becomes less successful. Thus, our

present understanding of the atmospheres of

pulsating variables, magnetic stars, WolfRayet stars, etc., is rudimentary at best.

Nevertheless, the tendency is to make quantitative measurements whenever possible and not

to be satisfied with visual estimates.

It is obviously impossible to cover all aspects

of such a varied subject in a short survey.

Fortunately several detailed discussions of solar

and stellar atmospheres have recently appeared,

and references to individual papers will be

found there.

References

AIXER, L. H.

1953. Astrophysics. Ronald Press, New York.

BEER, A., BD.

1955. Vistas in astronomy. Pergamon Press,

London and New York.

ClIANDRASEKHAR, S.

1950. Radiative transfer. Oxford University

Press, London.

HYNEK, J. A., ED.

1951. Astrophysics. McGraw-Hill, New York.

KoURGANOFF, V., AND BUSBRIDGE, I. W.

1952. Basic methods in transfer problems. Oxford University Press, London.

KUIPER, G. P., ED.

1953. The sun. University of Chicago Press.

UNSOLD, A.

1955. Physik der Sternatmospharen. Revised

ed. Springer, Berlin.

WOOLLEY, R. V. D. R., AND STIBBS, D . W. N.

1953. The outer layers of a star. Clarendon

Press, Oxford.

WRUBEL, M. H., ED.

1954. Proceedings of the National Science Foundation conference on stellar atmospheres.

Indiana University. Bloomington, Ind.

Aerodynamic Problems in Stellar Atmospheres

By Richard N. Thomas*

The term "cross-field research" occurs with

increasing frequency in astronomical studies.

This new emphasis does not imply that astronomy, in the past, failed to utilize advances

in other sciences; rather, it indicates our growing awareness that the physical problems of

matter studied by the observational astronomer are not far removed from those studied by

the laboratory scientist. Classically, astrophysics has dealt with matter under conditions that were strange and abnormal, compared to those in the terrestrial laboratory.

Astronomy has been an observational science.

The objects of its inquiry have been obscured

by the earth's atmosphere and situated at great

distances; and the phenomena studied were

incapable of reproduction in the laboratory.

Of all the "cross-field" liaisons that exemplify

our changing views, none is more interesting

than that between modern aerodynamics and

astrophysics.

Modern astrophysics not only concerns itself

with the common problems in the two sciences

but also recognizes that the data required to

support a theory or an anticipated future

development in one field may already exist in

the other. Also, the analytic methods that

are gradually being developed for a specific

problem in one field may already be common

knowledge in the other. For example, the

advent of high-altitude rockets, satellites, and

the possibility of extraterrestrial manned objects will break the barriers imposed by the

terrestrial atmosphere and the astronomical

distance scale. Yet a satisfactory development

of ultraspeed aerodynamics requires careful

attention to the problems of ablation from a

solid moving through a gas, of radiative energy

loss from a gas compressed rapidly to a high

degree, of the attendant ionization and chemical

reactions in the gaseous atmosphere, and of the

1

 Harvard College Observatory, Cambridge, Mass.

methods of description of a gas under conditions

departing drastically from local thermodynamic

equilibrium.

These problems are all of vital interest to

astrophysics because our observations show

that such configurations occur in astronomical

objects. Several of these problems, those connected with the ablation from a rapidly moving

solid object, form the fundamental basis of

meteor astronomy and have led to the creation

of astroballistics, a new field comprising both

astronomy and aerodynamics, which is described elsewhere in this series of articles.

Others of these problems seem to occur with

increasing frequency as we acquire more and

better data on stellar atmospheres and are

thus faced with the need to construct a theory

sufficiently precise to explain the data. In

this review, we want to elaborate on certain

aspects of stellar atmospheres and to emphasize

the continued interplay of astrophysical theory

and observation with aerodynamic theory and

experimentation, in a kind of checkerboard

pattern. The exciting aspect concerns the

common physical conditions occurring in the

stellar atmosphere and the aerodynamic experiment. The frustrating aspect is the lack of

facilities for experimental astrophysics, as it

becomes more and more evident that we can

actually duplicate in the laboratory physical

situations of direct astronomical interest.

Ten years ago, theorists concerned with

stellar atmospheres were completely preoccupiod with the notion of a gaseous atmosphere

whose properties were fixed by the three

assumptions: radiative transfer of energy, local

thermodynamic equilibrium, and the absence

of any velocity fields large enough to affect

the structure of the atmosphere. The discovery

by Schwarzschild, 50 years ago, of the overwhelming efficiency of radiative energy transport, relative to convective, completely dom123
124 SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS VOL.1

inated all thinking in subsequent years. The

later recognition of atmospheric convective

zones in stars where hydrogen was not mainly

ionized was of minor consequence; even in

these zones radiation dominates the energy

transfer. There were, however, three "types"

of atmosphere where "contrary" evidence was

becoming difficult to rationalize away: the

atmospheres of the variable stars, where the

notion of a velocity field had long existed but

where the exploration of the physical consequence of its existence had been rudimentary;

the atmospheres of certain supergiants, where

the evidence for large velocity fields had existed

for 10 years but nothing had been done in the

way of incorporating them into an atmospheric

model; and the solar chromosphere, where

considerable evidence of velocity fields and of

nonequilibrium phenomena caused the breakdown of virtually all our classical notions of a

stellar atmosphere.

There are two ways in which a velocity field

can perturb an atmosphere: through a transfer

of momentum or a transfer of energy. What

little thinking on the influence of velocity fields

that did take place in these earlier days was

concerned largely with the momentum effect.

Velocity fields were simply imposed as an added

pressure term, and the thermal and radiative

properties of the atmosphere were supposed to

remain invariant. Unfortunately, the velocity

fields discussed are often of such size that, if the

temperature of those atmospheric regions where

the velocity fields exist are not higher than we

would infer from our customary theories, the

motions must be supersonic. Such a situation

made it difficult to introduce the experience of

the aerodynamicist in describing velocity fields

in a compressible gas because such supersonic

motion, particularly of the quasirandom sort required by the astronomical data, is in aerodynamics always accompanied by Shockwaves and

by the rapid conversion of mass motion into

thermal. Thus we had the unhappy situation of

an empirical fact (large velocities) forced into an

artificial framework (no dissipation of mechanical energy) because astronomers were reluctant to abandon the atmospheric models

which had hitherto worked so well. Moreover,

there were absent many of the features one

would expect to occur in a rapid production of

thermal energy—emission lines, excesses of

ultraviolet radiation, etc. The situation, therefore, was far from clear cut.

(If the reader expects the above history to be

climaxed by a happy solution in this review, he

will be disappointed. If these problems had

been solved there would be little point in recording their apparently illogical history. We dwell

on the above pattern of thinking because, unhappily, much of it is today's pattern of thinking also. At present we can report only on

efforts to resolve these difficulties, not on their

success.)

On the other hand, when the thermal dissipation of the mechanical energy of the velocity

field was included, the effect on the atmospheric

model was, and is, likewise far from clear. This

energy transfer represents the most serious and

the most fascinating problem of all. First, any

such energy transfer establishes a cyclic process,

a mechanical energy source and a radiative sink.

Such a situation is ideal for setting up a departure from the conditions of local thennodynamic equilibrium. Second, such energy transfer

introduces the possibility of a perturbation on

the radiation field produced by the star. Consequently, the obvious subject for investigation

is the problem of the energy transfer associated

with velocity fields in stellar atmospheres.

The solar chromosphere offers the most natural starting point for an investigation of this

sort. First, from a general point of view, there

is one strong boundary condition on this investigation of aerodynamic effects in stellar atmospheres: the "classical" stellar atmosphere models have been remarkably successful in reproducing the features of the continuous spectrum

in the visual region. How can we introduce

such things as Shockwaves without affecting the

continuous spectrum? In the sun, the large

velocity fields seem confined to the outer atmosphere, where the total mass is too small to

make appreciable contribution to the continuum. We may, however, determine how far

down toward the photosphere the effect of the

chromospheric velocity fields extends by looking at weaker and weaker lines and the wings of

the strong lines. With this information, we may

be able to answer the question of what observable consequence would result from an increase

by some factor in the energy of the velocity field.
NEW HORIZONS IN ASTRONOMY 125

Or, alternatively, what would we have to do to

turn the solar atmosphere into that of a WolfRayet star, where the whole atmosphere seems

a composite of aerodynamic motions?

Second, with a view to finding the origin of

these velocity fields, the closeness of the sun

and our ability to resolve its disk are of enormous value. Many astronomers believe that

the chromospheric velocity fields simply reflect

the motion of the solar spicules, which are

supposed to arise from the granulation, whose

origin lies in the hydrogen convection zone.

Some of these astronomers regard magnetohydrodynamic effects as essential in the production of spicules, from granules; others believe

the transition can be effected from the aerodynamic behavior of a compressible gas in a

gravitational field alone; and still others believe

that the connection between spicules and

granules has yet to be established. Yet in all

these points of view several experimental and

theoretical aerodynamic problems stand out—

problems which we can investigate jointly in

astronomy and in aerodynamics: (1) What is

the behavior of a supersonic jet when radiation

processes form a major dissipative mechanism ?

Recent solar work on the stability of an

ionized gas against radiative cooling suggests

that these stability considerations may fix, to

a major degree, the permitted kinetic temperature of the gas in the jet. (2) How does a

supersonic jet composed of an ionized gas behave in a magnetic field, with the field (a)

parallel to the jet axis, and (b) perpendicular?

(3) Suppose we have a velocity field, with timedependent terms, in a gaseous atmosphere in a

gravitational field. How rapidly do the acoustic waves, produced by this field and propagated upward, die out? What is the interference pattern ? (4) If we superpose a fluctuating magnetic field on this velocity field with

time-dependent terms in an ionized medium,

can induction effects accelerate a nontrivial

amount of gas? (5) Last but hardly least in

our incomplete list of problems, How do we

specify the spectroscopic state of the gas under

the environment of these four problems, and

others in the same physical sequence ?

Let us return for a moment to the suggestion

that spicules are the source of the mechanical

energy which causes the chromosphere and that

381014—66 10

spicules arise from granulation. In this case,

all stars with atmospheric convection zones

should have chromospheres, that is, all mainsequence stars of spectral class later than~^40.

But what of the earlier spectral classes?

Should we expect no extended atmospheres

here, no velocity fields? Yet the existence of

differential velocities between emission and

absorption lines in stars of classes O and B

is well known. In a more extreme case, we

refer again to the Wolf-Rayet stars, where the

whole atmosphere seems composed of a system

of jet-ejections, or streams of outgoing and ingoing material. Whence arises this motion?

Aside from the possibilities of the hydrogen

convection zone (and possibly a helium convection zone in the hotter stars ?) and induction

effects in time-dependent magnetic fields, we

have two other sources of mechanical energy:

stellar rotation and stellar pulsation.

It is interesting to note that the first suggestion that aerodynamic turbulence might arise

in stellar atmospheres arose from the observation of the high Reynolds numbers associated

with stellar rotations. Furthermore, we note

that "shell" stars are most often those having a

rapidly rotating nucleus. What is the aerodynamic state of such a rapidly rotating gaseous

atmosphere ? Does the star set up a system of

aerodynamic turbulence, anisotropic for the

largest scale motion, with the greatest velocities

along the direction of rotation? This would

conform to laboratory experience. Or, is the

rotation velocity so high (greater than our

laboratory experience) that rings of matter are

ejected? This problem has hardly been

touched, astronomically or otherwise.

The concept of stellar pulsation is relatively

old; the implications of stellar pulsation,

relatively new. Given a gaseous atmosphere

of arbitrary temperature and density distribution, and an impressed oscillation of some thermodynamic variable at some level in the

atmosphere, what happens? Remember, we

are not dealing with the terrestrial atmosphere,

with temperature ~ zero so far as activation

of internal degrees of freedom are concerned. We need a terrestrial disturbance

of ~Mach 10 to cause any but the most trivial

luminosity. In the star, the ambient temperature ranges upward from~4,000°, and even
126 SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS

for Mach 2 one expects to find an observable

change in ionization and excitation. A literal

interpretation of velocity-curves would suggest

Af>2 in a wide class of variable stars. Indeed, in some stars emission lines appear over

some portions of the cycle, and we might

interpret their occurrence as implying the

existence of localized regions of high kinetic

temperature. On the other hand, some stars,

whose velocity variations imply compression

waves of sufficient amplitude to be shockwaves, do not show these emission lines.

Should we interpret these differing results as

suggesting that the compression wave just

does not have a sufficiently long path to develop

into a shock before it is out of the visible

atmosphere?

Actually, two physical parameters seem to

play a major role in aerodynamic phenomena

in the atmospheres of variable stars: the extent

of the h3Tdrogen convection zone and the extent

of the atmosphere. The convection zone

appears to determine the phase relation between

the luminous flux and pressure at the "origin"

of the disturbance; the extent of the atmosphere

determines the character of the disturbance,

and, in a sense, has a self-amplifying aspect.

The more the dissipation of energy in the atmosphere, the more extended the atmosphere,

hence the greater the chance for a shock to

develop and dissipate energy.

All these problems are tantalizing. Few

have been investigated, none solved. Our

greatest progress has been in our finally recognizing what the problems actually are. Our

greatest challenge is the need to develop some

physical insight to the solution of these problems which lie on the most advanced frontiers

of both astrophysics and aerodynamics.
Emission Nebulae

By Lawrence H. Aller

 1

Galactic nebulae that show an emission line

spectrum are of two types. Those associated

with the type I population are the so-called

diffuse nebulae; those associated with the

type II population are the planetaries. The

diffuse nebulae are normally irregular in form,

often of low density and surface brightness,

and sometimes of considerable size (e. g.,

Orion, 30 Doradus in the Large Magellanic

Cloud, and NGC 604 in the Triangulum

Spiral). The planetaries are often symmetrical

in form and smaller than the diffuse nebulae.

Many have a higher surface brightness and

higher density than the better known diffuse

nebulae.

Types of observations

Because of their different characteristics, the

kind of equipment and technique best suited

for the one type of emission nebulae is not

necessarily best for the others. In studies of

diffuse nebulae, wide-angle cameras of high

speed and detectors sensitive to low surface

brightnesses are desirable. For studies of the

planetaries a large scale (and therefore a large

telescope) is essential, since the angular diameters of all but the largest of these objects

amount to only a few seconds of arc. The

principal types of observation that can be

made are the following:

(a) Direct photographs with the aid of narrow band-pass filters to isolate monochromatic

nebular radiations. In particular, interference

filters have proved very useful in picking up

faint Ha emission regions throughout the

Milky Way. When most of the nebular radiation is concentrated in a few strong, wellseparated lines, the filter technique is extremely

valuable for studying problems of stratification

and structure.

1

 University of Michigan Observatory, Ann Arbor, Micb.

(b) Photographs with the aid of FabryPerot etalons and suitable plate-filter combinations are extremely valuable for studying

the internal motions in extended diffuse nebulae. This technique, introduced by Fabry

and Buisson more than 40 years ago and

further developed by Baade, Goos, Koch, and

Minkowski, has been applied with great success

by Georges Courtes at Haute-Provence. Studies may be made with relatively small wideangle telescopes that will yield considerable

information on the mass motions and internal

turbulence. It is important that such investigations be carried out not only in Ha but

also in the forbidden radiations of [O II] and

[O III] whenever possible.

(c) Spectrograpbic studies fall essentially

into three categories. First are slit spectrographic studies of the line and continuous

spectra to obtain line identifications and intensities from which the physical state in the

nebula or nebular filament may be inferred.

For this type of problem a so-called nebular

spectrograph with high speed and low dispersion is necessary for all but the brightest

objects. Planetary nebulae have been studied

rather intensively, particularly by Bowen and

his coworkers, but the spectra of diffuse nebulae

(except Orion and a few special objects) have

not been adequately investigated. Spectrographs of high dispersion may be used for

studies of the internal motions in small, bright

planetaries. Olin Wilson's multislit has proved

particularly valuable in such studies. Finally,

slitless spectrograms are essential for the study

of monochromatic images of small nebulae

when radiations that fall close to one another

are to be separated; e. g., the [N II] X6548,

6584 and Ha 6563 Lines.

(d) Polarization measurements are important

in these objects where scattering by small solid

127
128 SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS VOL. 1

particles occurs and in radio sources. Electron

scattering is rarely a factor in the continuous

spectra of diffuse nebulae.

(e) Radio-frequency observations of the freefree emission from diffuse nebulae have been

made by Haddock and his coworkers. Eventually, with the development of larger antennae,

it should be possible to measure the emission

from the brighter planetaries as well. These

objects have high surface brightnesses but small

areas so that their total flux is small. The

21-cm data supply information on those regions

of space where hydrogen is neutral.

Observable parameters

The following data can be obtained for typical

emission nebulae:

(a) The apparent size, form, and structure

Can be measured in each of the stronger monochromatic radiations.

(b) The surface brightness and isophotic contours can be measured for each of the stronger

monochromatic emissions.

(c) The wavelengths, identifications, and intensities of the lines and the energy distribution

in the continuous spectra can be measured for

typical nebulae.

(d) The distances of diffuse nebulae may be

found from the apparent magnitudes, color

excesses, and luminosity classes of their illuminating stars. The distances of the planetaries

eannot be established so easily, and for the

moment we shall have to be content with

statistical distances based on the association of

these objects with the central bulge of our

galaxy and external galaxies.

(e) The relationship between the structure

and internal motions can be established from a

comparison of isophotes of slitless spectrograms

or monochromatic photographs and radial velocity measurements. It is possible to establish

whether the motions are orderly expansions or

turbulent streaming, and to correlate the motion

with the appearance of the object. O. C.

Wilson has given a rather complete picture for

the planetaries. Similar studies for the diffuse

nebulae should prove worthwhile.

Interpretation of the observational data in

terms of physical parameters

The densities and temperatures of the radiating

gases may be established from the surface

brightness and from spectroscopic data. If the

nebula is smooth and unobscured and an

estimate of the electron temperature is available, a measurement of the surface brightness

in Ha or preferably in the Balmer continuum

will provide an estimate of the electron density.

Both diffuse nebulae and planetaries are frequently highly obscured. If a measurement of

the Balmer decrement is available, an estimate

of the space absorption can be found readily.

Albert Boggess III has shown how one may

combine optical measurements of surface brightness in Ha with radio-frequency measurements

to get the space absorption and densities for

certain diffuse nebulae for which Balmer decrements have not been measured.

The forbidden lines provide estimates of the

electron density N« and give us our best guesses

concerning the electron temperatures T(. If all

the radiations originate in the same strata one

may obtain N( and T« from a comparison of the

forbidden line intensities of [O III] and [N II],

for example. Often, however, the [N II] and

[O III] lines are predominantly produced in

different regions of the nebula. Hence, some

other method is needed. A number of years

ago, Ufford, Van Vleck, and I called attention

to the fact that the intensity ratio of the 3726

and 3729 [O II] lines varied from one nebula to

another, apparently depending on the density.

Seaton calculated improved target areas for the

collisional excitation of the [O II] lines so that

quantitative estimates of the density could be

made from the observed line intensities. Osterbrock has recently derived densities in the filaments of the Network Nebula, in the Orion

Nebula, and in other objects from the X3726/

X3729 data.

Some outstanding problems

At the moment, the physical nature of the

excitation of the bright line spectra of the gaseous nebulae appears to be well understood. The
NEW HORIZONS IN ASTRONOMY 129

continuous spectra of normal, purely gaseous

nebulae are less well understood, but seem to be

interpretable in terms of recombination of hydrogen and helium ions with electrons and the

double-photon emission. Further quantitative

measurements of the nebular continua are

needed to test various theoretical suggestions.

Further calculations of transition probabilities

for forbidden lines, especially those of iron in

various stages of ionization, are needed. Continuous absorption coefficients should be calculated for abundant ions, and further target

areas for the collisional excitation of metastable

levels should be computed. The accuracy of

the available theoretical target areas ought to

be improved.

Numerous theoretical problems connected

with the structure and especially the kinematics

of gaseous nebulae remain to be formulated and

solved. For example, what holds together the

thin filaments in such an object as the Network

Nebula? It is generally assumed that magnetic

fields are responsible, but the character and particularly the origins of these fields remain obscure. Any real understanding of the forms of

diffuse nebulae must await some fundamental

progress in the aerodynamics of compressible

ionized gases moving at supersonic velocities in

the presence of magnetic fields. Classical hydrodynamics and turbulence theories will require

much amplification and extension before they

can be applied to the diffuse nebulae and

planetaries.

The relation between the diffuse gaseous

nebulae and the solid particles of the interstellar medium requires further theoretical and

observational study.

Perhaps the most exciting problems are those

posed by the gaseous nebulae that are strong

radio sources. Oort and Walraven recently interpreted the continuous spectrum of the Crab

Nebula as "synchrotron" radiation from fast

electrons moving in a magnetic field. Possibly

such radiation is always or nearly always associated with strong radio-frequency emission,

although powerful radio sources do not necessarily emit a strong optical continuum. The

origin of the magnetic fields and the source of

the energy of the particles is not explained. It

is not sufficient to blame it on the supernova

process; the detailed mechanism must be worked

out.

New equipment and observational techniques

have expanded and are continuing to expand

our empirical knowledge of emission nebulae.

Theoretical interpretations of the forms and

kinematics are lagging behind. The problems

posed are among the most difficult in theoretical

astrophysics, a fact which makes the subject of

bright-line nebulae all the more challenging.

Some suggested references

The following books and review articles contain further references to the literature:

ALLER, L. H.

1956. Gaseous nebulae. Chapman and Hall,

London.

VAN DB HULST, H. C , BD.

1955. Second symposium on cosmical aerodynamics held at Cambridge, 1953. Published

for Int. Astron. Union by North Holland

Publishing Co. and by Interscience Press,

New York.

WURM, KARL

1951. Die planetarische Nebelu. Akademische

Verlag, Berlin.

Double and Multiple Stars

By Otto Struvex

The emphasis at the present time is upon

problems connected with the origin and evolution of double stars. G. P. Kuiper's discussion

in three articles in the Publications of the Astronomical Society of the Pacific2

 in 1935 and 1955

forms a sound basis for future work. He attributes the origin of double stars to the formation

of separate condensations in a primordial interstellar cloud which possessed too much angular

momentum, or too much mass, for the formation of a single star. In this respect the origin

of a multiple system does not differ significantly

from that of a star cluster. But the later evolution of a binary differs from the evolution of a

loose cluster. A visual double star is a very

stable formation and is not likely to change

very greatly during many hundreds of millions

of years, while a galactic star cluster may be

dispersed in the Milky Way in an interval of ten

or a hundred million years (though at least one

cluster, NGC 67, is believed by H. Johnson and

A. Sandage to have an age of 5X10* years).

Since we have good reasons for believing that

some galactic clusters are very young (the associations of Orion and f Persei, the double cluster

in Perseus, NGC 6231, etc.) while others are

much older (Pleiades, Hyades, Praesepe, NGC

725 and NGC 67), it would be of great interest

to know whether any binaries in these clusters

show differences in their properties that might

be related to the differences in their ages. As

yet, our information concerning all types of

binaries in galactic clusters is too fragmentary

for us to draw any conclusions.

Of special interest is the question of the evolution of close binary systems. From Kuiper's

work on /3 Lyrae, we know that two types of

instability may occur, with matter escaping

from one or both components either through the

1

 Department of Astronomy, University of California at Berkeley,

Calif.

» Vol. 47, pp. 15, 121, 1935; vol. 67, p. 887. 1955

bottleneck near the inner Lagrangian point, Llt

of the critical zero-velocity surface (instability

of type A), or through an opening at the first

outer Lagrangian point, Zj (instability of type

B). Kuiper thought that the instability of fi

Lyrae is a combination of types A and B, but

more recent work seems to indicate that both

components are unstable and lose matter in the

form of two streams moving in opposite directions through the bottleneck near L\. There

can be little doubt that this phenomenon produces the secular change in the period of 0

Lyrae. The loss of matter in /3 Lyrae, if the

streams are ultimately expelled into space, must

amount to about 5X10" gr/sec in order to explain the increase of the period. Since the mass

of the star is of the order of 103* grams, the life

expectancy of /3 Lyrae should be of the order of

100,000 years.

No other star is known that closely resembles

/3 Lyrae (in itself a good indicator of the short

duration of its particular stage of evolution).

However, many other systems have gaseous

rings or streams and many of them, according

to F. B. Wood, have changing periods, especially

when at least one component is a large subgiant

whose volume fills its loop of the critical zerovelocity surface.

One of the most pressing needs in this field is

a systematic, prolonged study of the periods of

all suitable eclipsing and spectroscopic binaries.

The photometric observations would be relatively simple, but should be made with photoelectric means. Results of great value would

become available in a few years; but even more

important would be a list of accurate periods

during the next 10 years which could later be

compared with similar lists at intervals of about

25 years. We cannot start too soon with such

a program!

The spectroscopic observations could, for

many systems, be carried out with telescopes

131
132 SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS

of moderate size. We need periodic determinations of the velocity curves of all the

brighter spectroscopic binaries. Here again

results may be expected immediately, not only

with regard to the periods but also with regard

to changes in the orbital elements. Algol,

Spica, 5 Orionis, and many more should be

reobserved at least once every 10 years.

One of the most puzzling questions in the

field of double-lined spectroscopic binaries is

the interpretation of the tendency of many

systems to show a strengthening of the violetdisplaced absorption lines relative to the reddisplaced components, irrespective of which

of the two stars happens to be approaching.

This phenomenon was discovered by Miss

Maury at Harvard in the spectra of n Scorpii

and V Puppis. But later observations at

McDonald Observatory failed to show this

variation. Yet, independently, we found that

it was conspicuous in Spica and in several other

short-period spectroscopic binaries of early

class B or O. A recent study of v Scorpii by

S. J. Inglis shows that it is present in the

helium lines but that it may be reversed in

sign in the hydrogen lines. I had previously

assumed that there is more gas in front of

the advancing hemispheres of each component

than in front of the receding hemispheres.

However Inglis' result suggests that it may

be caused by a difference in ionization and

excitation. As far as I know, there are no

photoelectric measurements that would help

us in explaining this phenomenon. Since it

occurs only in close systems (it is also conspicuous in many W Ursae Majoris systems)

I suspect that it is in some way related to the

existence of gaseous streams.

S. S. Huang has recently asked whether

there might not be accumulations of gaseous

material at those two Lagrangian points which

correspond to the equilateral triangle solutions

of the restricted three-body problem. Unless

the mass ratio of the binary is very different

from one, these solutions are not stable, but

there could still be regions of increased density

in the gaseous ring at L% and L4. It is possible

that a more careful study of the eclipses of the

emission features in the hydrogen lines of RW

Tauri, RW Persei, SX Cassiopeiae, U Sagittae,

and many more would enable us to locate the

regions in which these emission lines are

produced. Joy's original interpretation of RW

Tauri in terms of a ring around the brighter

component of the binary, or my later interpretation in terms of a ring of gas surrounding the

entire system, may have to be modified in the

light of Huang's idea.

During the interval 1940-1950 several

McDonald astronomers made a superficial survey of the velocity curves and spectroscopic

features of about 100 eclipsing variables. Most

of the systems presented no peculiarities and

the results were sufficient for statistical studies,

but a considerable number yielded results of

exceptional interest. A few of these stars

have since been investigated more exhaustively

(U Cephei, U Sagittae, SX Cassiopeiae, RX

Cassiopeiae, RZ Scuti), but quite a number

have not yet been scrutinized. I can mention

here only a few of them. WX Cephei (J.

Sahade and C. U. Cesco) shows amazing

changes in the intensities of the two Ca II

absorption components. They are unlike the

changes discussed above. DN Orionis has an

exceptionally small range of velocity and,

therefore, a very small mass function. It is

one of the most peculiar binaries in this respect,

resembling XZ Sagittarii (J. Sahade). Both

must possess subgiant components of very small

mass (about 0.1 or 0.2©) but of large diameter. These components violate the conventional mass-luminosity relation. AW Pegasi

and V 377 Centauri (Sahade) are even more

peculiar, and we have been unable to interpret

their spectra. Details about these and other

systems may be found in the Astroph3rsical

Journal or in the Contributions from the

W. J. McDonald Observatory.

The famous visual double star 7 Virginis

has been suspected by Lick astronomers to

vary in light. Since both stars can be observed

spectroscopically with high dispersion, a test

of each component (or of the integrated light

of both) for variability would be of interest.

Any visual binary whose orbit is known should

be investigated at least in integrated light for

variability. The discovery of an RR Lyrae

or another pulsating component in a double star,

whose masses are known, would be a major
NEW HORIZONS IN ASTRONOMY 133

triumph. It might even be possible from the

colors (or spectra) and magnitudes to pick out

promising pairs, but the chances for success

can hardly be estimated. I recommend this

important search only to those who have

other problems that are certain to give significant results.

A few eclipsing binaries must be observed as

nearly continuously as is possible, both photoelectrically and spectroscopically. Among

these I should list e Aurigae, W Cephei,

/S Lyrae, UX Monocerotis (R. Lynds suspects

that the small, bright, and hot component of

this system is a short-period pulsating variable,

while the other component is a large and cool

subgiant). No spectrographic equipment is

now available for such continuing observations.

I therefore endorse Th. Dunham's project for

the construction of a new spectrographic

telescope.

I hope that someone will have observed the

light curve of « Aurigae over a wide range of

wavelengths. It is important to know whether

the depth of the eclipse (a totality began Dec.

15, 1955, ending 330 days later) is really the

same in all wavelengths.

What is the dispersion of the mass-luminosity

relation, expecially at the upper end? If

nuclear evolution proceeds without appreciable

change in mass, we should expect some scatter

among the O- and B-type stars. Most of our

present information comes from the velocity

curves of double-lined binaries, and this problem is looked after by the astronomers at Victoria. Special attention should, however, be

devoted to such stars as J. S. Plaskett's massive

binary, in which the more massive component

has the weaker absorption lines; the profiles

of the lines of this component show irregular

changes.

It is not yet clear whether the yellow and red

giants in Population I obey the conventional

mass-luminosity relation. K. O. Wright's work

on Capella seems to indicate that it does, and

my own earlier measurements suggested only

a relatively small departure. D. M. Popper's

work on f Aurigae and other similar systems

may soon solve this problem. It may help

us in constructing the evolutionary tracks of

massive stars, but here again we must be

prepared to find that the giants lose mass by

expansion. A. J. Deutsch's work on the expanding shell of the system of a Herculis supports this hypothesis.

What is the nature of the so-called Trumpler

stars? Some of them occur in binary systems

but we have as yet no firm knowledge of their

masses. As Trumpler has shown, their lines

have large red shifts which, if interpreted by

the theory of general relativity, would give

very large masses (or very small radii). There

are some reasons for believing that their masses

are not more than about 75o> But how do

we then explain the red shifts?

Are all post-novae and all SS Cygni stars

close binaries? And why are there so many

Wolf-Rayet stars which are components of

close binary systems? Are all of these objects

the products of the evolution of binaries (as

Huang suggests in a recent article in the Astronomical Journal8

) and are the RR Lyrae stars

the products of nuclear evolution of single

stars (or of wide visual binaries) uncomplicated

by the loss of mass resulting from the instability

of a close binary? It is not yet clear why loss

of mass into space (as distinguished from the

formation of streams within a binary, or the

exchange of mass between the components of

a binary) should be significantly different.

We must assume that once a stream of gas is

generated in a binary it is somehow expelled

from the entire system, and that the absence

of such streams in single, slowly rotating stars

impedes their loss of mass by "corpuscular

radiation."

It would be unreasonable to conclude this

report without stressing the importance of

continuing routine observations of all classes

of double and multiple stars. In preparing

his program, an investigator should not be

concerned solely with the solution of immediate

problems but should devote a part of his time

to the accumulation of accurate observations

which will bear fruit long after his death.

The most important astronomical problems—

those of evolution of stars, of perturbations in

double star orbits, of slow periodic or irregular

variations, and many more—require tens or

hundreds of years of persistent effort. It is

•VoLftl, p. 40, 1956.
134 SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS

especially difficult to sustain an enthusiastic every astronomer has profited from the observaattitude on the part of a scientist if he does tions of his predecessors; he should therefore

not expect to gather the harvest resulting be willing to pass on to his successors a store

from his labor. Nevertheless, we must all of good observations to make up for what he

follow the dictum of Ejnar Hertzsprung: has inherited from former generations.
Stellar Dynamics

ByU. Van Wijk1

Although stellar dynamics is largely a theoretical field, anyone working in it has a strong

interest in observational results concerning the

motions of the stars—their radial velocities and

proper motions.

One of the more active areas of stellar dynamics has been the study of the interaction

between stars and interstellar clouds insofar as

it affects the motions of the stars. In papers

by Spitzer and Schwarzschild (1953) and by

Osterbrock (1952), the effect of randomly distributed interstellar clouds has been considered.

One may hope that the 21-centimeter band

observations will ultimately lead to a more

specific idea about the distribution of interstellar clouds and that the effect of nonrandom

distributions, such as a spiral arm, may then

be considered.

A field which is basically quite fruitful but

which has been lying fallow recently is that of

spherical clusters and, particularly, the escape

of cluster stars. The papers of Spitzer (1940)

and Chandrasekhar (1943) are still the basic

references, and it may well be that further

progress will strongly depend on the ready

availability of electronic calculators.

An interesting recent development is the

1

 Princeton University Observatory, Princeton, N. J

attempt to interpret the present kinetic properties of the stars in terms of the conditions

under which they were formed. Examples of

this approach are the papers by Belzer, Gamow,

and Keller (1951) and van Wijk (1956). Here

we approach the border line between stellar

dynamics and the study of galactic structure.

But two subjects which are just beyond this

border line may still be mentioned here because

of their intimate connection with our present

subject: the work on pseudointegrals of motion

by Lindblad and the study by Blaauw (1952)

of the dissolution of very loose galactic clusters.

References

BELZER, J., GAMOW, G., AND KELLER, G.

1951. Astrophys. Journ., vol. 113, p. 166.

BLAAUW, A.

1952. Bull. Astron. Inst. Netherlands, vol. 11,

p. 414.

CHANDRASEKHAR, S.

1943. Astrophys. Journ., vol. 98, p. 54.

OSTERBROCK, D.

1952. Astrophys. Journ., vol. 116, p. 164.

SPITZER, L., JR.

1940. Monthly Notices Roy. Astron. Soc. London,

vol. 100, p. 396.

SPITZER, L., JR., AND SCHWARZSCHILD, M.

1953. Astrophys. Journ., vol. 118, p. 106.

VAN WIJK, U.

1956. Astron. Journ., vol. 61, p. 277.

135

Variable Stars

By C. Payne-Gaposchkin*

Variable stars play a dual role in astronomy.

First, they furnish standards of luminosity and

distance. Second, their own physical properties throw light on energy sources, the processes

of pulsation and explosion, and the circumstances of stellar formation.

Standards of luminosity

The RR Lyrae stars, Cepheids, novae, and

supernovae have been the most important

indicators of distance. Their average luminosities (mean for the pulsating stars, maximum

for the explosive ones) are known with an uncertainty of a few lOths of a magnitude. We

are now concerned with determining the dispersion of luminosity for stars of these groups.

The absolute magnitudes of RR Lyrae stars

will emerge from critical comparisons of the

color-magnitude arrays of globular clusters.

Probably they will be found to be related to the

metal richness of the cluster stars. The

striking differences between the noncluster RR

Lyrae stars of our neighborhood and those near

the galactic center, when interpreted in the

light of the data from globular clusters, will

furnish a new approach to the dimensions of

our galaxy. These studies, which require

spectroscopic analysis and photoelectric accuracy, can be made only with the largest

telescopes. The extension of this work to the

nearby, bright globular clusters of the Southern

Hemisphere would be most important.

The dispersion in the absolute magnitudes

of Cepheid variables and its relation to color

and spectrum is likewise an investigation to

be carried out in the Southern Hemisphere.

The Cepheids in the Magellanic Clouds furnish

perhaps the most urgent and profitable problem

in applied variable star astronomy; the study

should be made concurrently with photoelectric and spectroscopic equipment. It should

> Harvard College Observatory, Cambridge, Mass.

be borne in mind, however, that the Magellanic

Cepheids may differ in detail from galactic

Cepheids of like period. The study of color

and spectral variation of galactic Cepheids

should therefore be pressed, with especial

attention to the effects of selective absorption.

Here, again, material from the Southern

Hemisphere is of particular urgency; the

rich Carina region seems comparatively free

of absorption, and radial velocities of Cepheids

in this area will fill out our information on the

relation of the Cepheids to galactic rotation,

and thus on the dimensions of the galactic

system.

The absolute magnitudes of novae have

recently been placed on a firm basis by the

light curves of the novae in Messier 31. Probably the novae in other nearby galaxies will

be the most valuable means of extending this

information.

The absolute magnitudes of supernovae,

the distinction between the two types, and their

relationship to ordinary novae must again

rest on studies made with the largest telescopes.

Physical problems

Detailed study of the brighter variable

stars is continually revealing new types of

variation. Future progress in variable star

astronomy calls for understanding the relationships of these types to one another and of the

physical processes involved in them. While

the discovery of variable stars should certainly

continue, it seems probable that systematic

study of chosen typical specimens is more

important than indiscriminate search for faint

variables. However, it should not be assumed

that all undiscovered variable stars are faint;

even among the naked-eye stars there may still

be a number of important examples hitherto

undetected. Means for tracing the variations

of brightness of stars that are spectroscopically

137
138 SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS

interesting, such as shell stars and Of stars,

should not fall into abeyance; in this regard

the maintenance of adequate sky patrols

(not necessarily with powerful instruments,

as the profitable stars are likely to be brighter

than the tenth magnitude) is a project worthy

of continued support in both hemispheres.

It is important that the variable stars chosen

for systematic investigation be studied with

the greatest available precision. Brightness

and color should be measured photoelectrically.

Spectra should be photometrically analyzed,

and all spectra should be suitably standardized

for this purpose. Estimates of the intensities

of absorption and emission lines are not good

enough for modern purposes. It should be

emphasized that the red and near infrared

are now as accessible as the photographic

region, and often provide data that are equally

or more significant.

Lastly, a star that is varying should be

studied concurrently by photometric and spectroscopic methods; this is of especial importance

when the variation is not strictly periodic.

Eclipsing stars.—Of unique importance because of their contribution to our knowledge of

stellar masses and dimensions, eclipsing stars

have entered a new era with the coming of

photoelectric photometry. Any eclipsing star

that gives a two-spectrum velocity curve and a

good light curve adds to our fundamental

knowledge of stars.

Several groups of eclipsing stars give unusual

promise. The combinations of a blue star and

a red giant with an extensive atmosphere are

important in the study of stellar envelopes; in

addition to the well-known specimens now being

actively studied, others are undoubtedly to be

be found in lists of "composite spectra."

The remarkable analogies between the shortperiod UX Ursae Majoris, the system of Nova

Herculis, and the pair that constitutes AE

Aquarii invite a search for such stars, which

will throw light on the processes of stellar

explosions.

Periodic variables.—The classical Cepheids and

the parallel sequence of Population II variables

are rich in possibilities. Concurrent studies of

brightness, color, spectrum, and radial velocity

will be needed before the interrelationships of

these groups are clarified. Careful spectrophotometry will probably reveal differences of

composition. Equally important is the development of the physical theory that bears on the

propagation of waves through the envelopes of

these groups of stars.

The long-period and other red variables involve similar theoretical problems. Their

spectra also invite quantitative observational

work: the identification and intensities of

bright lines are of obvious importance, and

quantitative study of the absorption spectra

(admittedly beset with grave observational

difficulties) is also a problem of the first

importance.

Explosive variables.—The most difficult problem, which can be attacked only with the

largest telescopes, involves the physical processes of the supernovae, and especially the

unsolved problem of their spectroscopic changes.

The study of the novae proper has always

been hampered by the unexpectedness of the

outbursts. Too often, alas, a nova has been

studied from a sense of duty, and each contribution thus tends to be isolated from related

studies. Every suitably equipped observatory

should have a nova program prepared in advance, and the supplies necessary to put it into

action. The program should from the first

include the red and near infrared, and especial

attention should be paid to photometric

standardization.

The U Geminorum stars are so faint that the

necessary spectroscopic studies require powerful equipment. The likelihood that all are

double stars and the possibility that many are

spectroscopic binaries of short period make

them of exceptional interest. Photometric

studies might reveal that some of them are

eclipsing stars; in addition, it would be important to learn whether the recently observed

rapid changes of SS Cygni are a common feature

in the rest of the class.

Irregular variables.—The faint irregular variable stars in nebulosity are destined to play an

important role in ideas of the origin and development of stars. The discovery of these "T

associations" and the concurrent spectroscopic

and photometric study of their behavior are

extremely important provinces of variable star

astronomy.
Interstellar Matter

By Lyman Spitzer, Jr.1

The study of interstellar matter may be

divided into two main parts. Firstly, there is

the measurement of the interaction between

interstellar material and electromagnetic radiation. The objective of this research is to

determine, from the observed radiation reaching

the earth, the properties of the material

between the stars. While this study is primarily observational, the interpretation of the

observations in terms of the number of absorbing atoms or of the properties of the absorbing

solid particles may involve much theoretical

analysis.

Secondly, there is the physical analysis of

those processes that occur in interstellar space.

In this study the density, composition, and

spatial distribution of interstellar material are

obtained from the observational research described above, and the various physical processes that occur are deduced from the laws of

nature, as determined in terrestrial laboratories.

Thus, for example, the internal temperature

of the small solid particles or grains is computed theoretically, the interactions between

atoms and grains are analyzed, and the methods

by which interstellar material can condense

into a new star are deduced.

At present our information on material in

space is seriously limited both by observational

inaccuracies and by a lack of information on

the relevant laws of nature. During the next

few decades significant advances in both these

respects may be anticipated.

New experimental techniques promise a

great increase in observational accuracy. Of

particular importance is the imminent development of photoelectric spectrophotometry. The

combination of a photoelectric cell and a

spectrograph will have many important advantages in interstellar research. One possibility, which is already being explored by

1

 Princeton University Observatory, Princeton, N. J.

Stromgren at the Yerkes Observatory, is the

accurate color measurement for stars of all

spectral types. Accurate measures of absorption lines yield very precise determinations of

spectral type, and hence of the intrinsic stellar

color. The color excess, or deviation between

the actual color and the intrinsic color, may

then be found for any star. With this development we should be able to determine the

density of small solid particles at points in the

sky very close together, and thus obtain a more

detailed picture than is now possible of the

distribution of interstellar material.

Another suitable program for photoelectric

spectrophotometry is the precise measure of

interstellar line profiles. With this technique

it is possible to determine both the number of

absorbing sodium and calcium atoms in the

line of sight to each star and the precise velocity

distribution of the atoms. Still another possibility is the measurement of much fainter

interstellar lines than are customarily observed.

Thus, detailed information can be obtained on

the spatial distribution of such interstellar

atoms as neutral Ca and ionized Fe and Ti.

At present the interstellar lines produced by

these atoms can only barely be detected in a

very few stars.

Radio astronomy is another very important

observational development in this field. As the

size of radio telescopes increases more and more,

the information about the distribution of neutral

hydrogen will become more and more detailed.

Information about other types of atoms may

also be obtained, perhaps, with radio techniques. Evidently our exploration of interstellar clouds will become much more precise

and much more extensive in the years to come.

This increasing detail of information, and

this very great increase in the accuracy of interstellar maps that can be drawn, will naturally

lead to an expanded effort on the analytical and

139
140 SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS VOL.1

theoretical front. For example, data are gradually accumulating on what might be called

interstellar meteorology—the great problem of

the motions and transformations of the interstellar clouds. As in terrestrial meteorology,

this subject must include not only ordinary

dynamics but also the effects of radiation, both

emitted and absorbed. Changes of state, involving the combination and dissociation of

molecules and atoms, must also be considered.

These phenomena are probably influenced profoundly by the presence of both large-scale

magnetic fields and cosmic rays, which have no

counterpart in terrestrial meteorology. Evidently the comprehensive analysis of the many

simultaneous processes will be a vast and challenging adventure.

Another theoretical problem concerns the

physical chemistry of grains. The boundary

conditions in this problem are simple. We

know approximately the conditions under which

atoms in interstellar space gather together to

form solid particles. However, the nature of

these particles is completely unknown. In particular, the detailed structure of the grains, the

forces holding them together, and, in fact, all

except the grossest properties of these small

solid particles are completely uncertain. A

precise, elaborate, and ingenious laboratory

approach, together with a considerable amount

of theory, is probably required to shed light on

the situation. The interactions of grains with

atoms, including the rate at which the grains

grow and are destroyed and the various selective ways in which they interact with different

types of elements, are of particular interest

astronomically. Another important problem

concerns the magnetic properties of such grains,

the solution of which also probably demands a

combination of observation and theory.

A central theoretical problem is, of course,

the one of evolution. How do clouds evolve

over a period of a billion years or more, and, in

particular, how are new stars formed from the

interstellar medium? The detailed answer to

these questions is perhaps the most difficult

part of the entire field, and from a logical standpoint should perhaps be postponed until the

rest of the subject is fully developed. In view

of the very great interest in this particular

problem, however, it appears certain that during the next few decades much new work will

be done. While no final results may be expected in this area for some time, it seems very

likely that ingenuity and hard work will produce

many glittering and intriguing hypotheses.
21-Centimeter Research

By Bart J. Bok '

The discovery of the 21-centimeter line of

neutral hydrogen by Ewen and Purcell (1951)

and the subsequent confirmation of this discovery by Dutch and Australian groups (Muller

and Oort, 1951; Pawsey, 1951) marked the

birth of a new era in the study of our Milky

Way system and of the universe of galaxies.

That radio radiation with a wavelength of 21

cm emanating from neutral atomic interstellar

hydrogen might be sufficiently strong to be

recorded was first suggested by van de Hulst

(1945) and was predicted independently by

Shklovsky (1949). But not until the actual

discovery of the line did astronomers begin to

realize how powerful a tool for future research

had become available.

Research in the 21-cm field may at present

be divided roughly into three broad areas.

The first concerns the spiral structure of our

galaxy, and has yielded spectacular first results.

Really detailed analysis, however, has only

begun, and will undoubtedly be fruitful in

future studies of spiral structure. The second

category is that of related optical and radio

studies of the fine structure of the interstellar

gaseous medium in our galaxy. We shall list

below some of the topics of research most

important here, but it is already clear that the

"21-cm astrophysicist" has waiting for him a

multitude of specific research problems. The

third category is that of extragalactic 21-cm

research, one in which the work has barely

started but which obviously will attain increasing importance as the reach of our equipment

increases.

One should bear in mind that up to the

present all but one of the five principal radio

telescopes that are active in the 21-cm field are

paraboloids with apertures of the order of 25

to 30 feet, the one exception being the 50-foot

1

 Harvard College Observatory, Cambridge, Mass., and Commonwealth Observatory, Canberra, Australia.

paraboloid at the Naval Research Laboratory

in Washington. Also, most of the work until

now has been done with electronic equipment

with bandwidths of 40 and 15 kc/sec (equivalent

to resolutions in radial velocity of 8 and 3

km/sec) and very few observations have been

made with bandwidths between 5 and 10

kc/sec (equivalent to radial velocity resolutions of 1 to 2 km/sec). Within a year or so

there should be in operation at least a half

dozen radio telescopes with apertures of 60 to

100 feet, all of them steerable and with surface

tolerances sufficiently strict to permit full use

at the 21-cm wavelength. Research at 21 centimeters has just been started with two large

radio telescopes—the 83-foot instrument at

Dwingelo in Holland and the 60-foot George R.

Agassiz radio telescope of Harvard Observatory

(see the recent article by Heeschen, in press).

Several of these instruments apparently will

have electronic equipment capable of recording 21-cm profiles with a bandwidth of 5 kc/sec

(equivalent to a resolution in radial velocity of

1 km/sec). The group at the California Institute of Technology is planning to use two 90-

foot antennas as an interferometer with variable

separation of the two component antennas.

Moreover, three very much larger paraboloids

intended in part for use in 21-cm research are

now either under construction or on the drawing board: the 250-foot paraboloid for Jodrell

Bank, the 140-foot paraboloid for the InterUniversity Radio Observatory now under consideration by the National Science Foundation,

and the large instrument planned for Australia.

Spiral structure of our galaxy

The Leiden group has contributed most

effectively to the study of the spiral structure

of our galaxy. The early work by the Leiden

group culminated in the famous joint paper by

van de Hulst, Muller, and Oort (1954) which

141
142 SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS

contains the picture of the spiral structure of

our galaxy in its broad features as observable

from a Northern Hemisphere radio observatory.

For the Southern Hemisphere a preliminary

survey has been completed by Christiansen

and Hindman (1952), and the most recent

results were presented by Carpenter (1955)

at the Troy meeting of the American Astronomical Society. Since the original Dutch

survey was completed, there has been considerable activity in the study of the spiral

structure for the part of our galaxy between

galactic longitudes 40° and 200°; that is, from

Cygnus to well beyond Sirius. The most

extensive research has been done by Westerhout

(in press), who has derived and analyzed 21-cm

profiles for the entire section between galactic

latitudes —10° and +10°. To supplement the

work of Westerhout, we have the published

and unpublished results of Matthews (1955)

and of Davis at Harvard and of Heifer and

Tatel (1955) at the Carnegie Institution. All

these researches indicate that the simplified

descriptions of three primary spiral arms in this

section, the Orion Arm, the Perseus Arm, and

the Distant Arm, are only very rough first approximations. The available evidence suggests major branching of the principal spiral

features; it is evident, furthermore, that all of

the spiral features are definitely not precisely

in the same plane. Comparable detailed studies

for the Southern Hemisphere are still lacking,

but the first results by the Australian groups

suggest that here, also, complexity rather than

simplicity are to be found.

The study of the spiral structure in the part

of our galaxy interior to our sun offers a variety

of complex problems. The principal source of

difficulty is that, even on the simple assumption

of purely circular galactic rotation, a feature

with a given radial velocity may be produced

by hydrogen clouds at either one or another

distance; or, more generally, that a particular

feature may well be the blurred combined

impression produced by features at two different

distances from the sun. The early results

obtained by the Leiden group have been

summarized in a paper by Kwee, Muller, and

Westerhout (1954); and, more recently, Schmidt

(in press) has obtained useful new evidence. At

Harvard, Heeschen and Howard (1956) have

begun a study rather similar in scope to that of

Schmidt. It seems established that one Inner

Arm is present which presumably is identical

with the optically discovered Sagittarius Arm,

but further evidence exists for the presence of

some structure in the internal neutral-hydrogen

conglomerate, close to the center of our galaxy.

Two major problems arise in any analysis of

spiral features from 21-cm profiles. First, a

large group of problems relates to the probable

shape of the curve that gives circular velocity

of rotation as a function of distance from the

galactic center. In 21-cm research we have no

direct way of obtaining distances, and we can

only assign to a given observed feature a distance based on its observed radial velocity of

approach or recession. If our assumed form

for the curve relating distance from the galactic

center to circular velocity of rotation is in

error, then all of our derived distances obviously will be in error also. A second and related

cause for difficulties lies in the fact that, on the

basis of 21-cm observations alone, we have no

way of disentangling effects of peculiar motions

from those caused by purely circular galactic

rotation.

We should be able to outline on a fairly

homogeneous basis the general features of the

spiral structure of our galaxy once the southern

Milky Way observations are completed and

analyzed as have been the profiles for the northern Milky Way by the Leiden observers.

Great interest attaches to the future clarification of the details of spiral structure, but here

progress will be slow at best, and, initially at

least, somewhat uncertain. We have already

indicated that the dynamical study of the law

of circular galactic rotation is imperative if

we wish to derive meaningful distances for the

features found through our 21-cm observations.

To deal conclusively with problems of the

branching of spiral arms and with fine structure

phenomena, our only hopeful approach appears

to be through coordinated optical and radio

studies of sections of the Milky Way, one by

one. For proper analysis of 21-cm features,

we shall need extensive and precise optical

data—for example, for OB stars and associations

in the area of the sky covered by a particular

21-cm study. For proper interpretation of the

21-cm features, we should at all times attempt
• a t NEW HORIZONS IN ASTRONOMY 143

to obtain radio profiles for a fairly wide range

of galactic latitudes and for points in the area

of investigation spaced as close together as

appears to be reasonable for a particular radio

telescope. Preliminary work in this field by

Matthews (1955) has already demonstrated

that the results of closely spaced surveys will

be well worth the effort, and that correlated

radio and optical studies will provide insight

into the fine details of spiral structure. Furthermore, a great deal of new information should

be obtainable regarding the fine structure of

the interstellar medium and the association

between interstellar hydrogen and certain

classes of stars. A related problem of great

interest is the relative distribution of neutral

hydrogen and cosmic dust inside a spiral arm,

or branch of an arm.

Astrophysical studies

The 21-cm astrophysicist is essentially a

spectroscopist who has for study only one

single spectrum line. Fortunately, the microwave spectrographs presently in use are really

high-dispersion instruments and thus much

can be learned from a single 21-cm profile.

The 21-cm astrophysicist realizes that he is

the fortunate fellow who has first achieved

access to the all-pervading medium of neutral

hydrogen of our galactic system, which means

probably as much as 95 percent of all the

interstellar hydrogen. In this section, we shall

describe first some astrophysical problems that

have proved practicable for paraboloid antennas with apertures of 20 to 30 feet, then turn

to problems whose solution is aided by antennas

with apertures of 50 to 80 feet, and conclude

with some astrophysical problems that we may

attempt to study with antennas of 100-foot

and larger apertures.

The work of Lilley (1955) and of Heeschen

(1955) at Agassiz Station illustrates the first

group—astrophysical problems that may be

studied with antennas of relatively small

aperture. For the Taurus Complex of dark

nebulosity, Lilley found generally far stronger

signals for fields well inside the dark complex

than for relatively unobscured fields bordering

upon the complex. It thus became evident

that the large dust complexes have excess

neutral hydrogen gas associated with them.

Lilley found the ratio between the average

densities of gas and of dust to be close to

100:1 inside the Taurus Complex. Lawrence,

Menon, and I (Bok et al., 1955) used the 24-

foot paraboloid to study some very densely

obscured spots inside the Ophiuchus and

Taurus Complexes. We found no marked

differences in 21-cm signal strength when we

compared intensities for the highly obscured

fields with those found for neighboring fields

still inside the complexes but definitely less

obscured. We deduced from our observations

that the large dust complexes are regions of

markedly high concentration of interstellar

atomic neutral hydrogen, but that the distribution of the atomic hydrogen inside the

complex does not mirror precisely that of the

cosmic dust. Existing 21-cm observations

suggest that the neutral hydrogen is more

evenly distributed than the cosmic dust, but

no really significant conclusions can be drawn

until the really large paraboloids become available for studies of this nature. A careful

search for neutral hydrogen concentrations by

21-cm interferometer techniques should prove

of great interest in connection with this problem; a negative result would be fully as important as a positive one. When the necessary

observations have been made, their interpretation will still present a good many difficulties.

Our lack of information regarding the presence

of molecular hydrogen complicates analysis

considerably, and as yet the possible effects

of variations in temperature inside the large

dust complexes are by no means understood.

The assumption that the interstellar medium

of neutral atomic hydrogen possesses one single

characteristic temperature is at best a very

crude one, and in future research temperature

variations undoubtedly will have to be considered along with density variations.

A second area in which 21-cm studies have

already made important contributions is that

of the study—by combined optical and radio

techniques—of regions of special interest, such

as the Orion region. The recent work by

Menon (in press) illustrates how much we can

learn through detailed 21-cm studies. Optically, the Orion region is known by its

separate features: the Great Nebula in Orion,

with a total estimated mass not in excess of
144 SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS

1,000 solar masses; the Trapezium Cluster and

various OB associations and aggregates;

Barnard's Luminous Arc of very great dimensions; and, finally, the large dust complex that

is evidently obscuring much of the region. The

21-cm observations by Menon show that all of

these features are imbedded in a large "bowl of

neutral hydrogen jelly," with a total mass

probably of the order of 60 to 100 thousand

solar masses. Menon's observations suggest

that this entire mass is in slow rotation and

that it is, very likely, expanding at the rate of

8 to 10 km/sec. Here the 21-cm approach has

already added a dimension to our earlier unrelated optical studies. The astrophysical and

cosmogonical importance of the new approach

is obvious.

The third group of problems that are proving

amenable to study by 21-cm techniques deals

with the fine structure of the interstellar

hydrogen medium. The studies by Lawrence

(in press) and by Heeschen and Drake (1956)

serve to illustrate two different approaches.

Lawrence, working in close collaboration with

Munch, has studied at the Agassiz Station

21-cm profiles for fields centered upon stars for

which Munch, at Mount Wilson and Palomar,

had found multiple optical interstellar lines.

Two established coincidences between radial

velocities of 21-cm features and strong interstellar absorption lines indicate that the same

interstellar gas clouds may be responsible.

Further study of these phenomena should prove

exceedingly rewarding, but we can hope to

arrive at a clear understanding through 21-cm

studies only with radio telescopes of large

apertures with associated electronic equipment

of very narrow bandwidth. The research of

Heeschen and Drake refers to the general

direction of the Pleiades Cluster. They find a

double-peaked 21-cm profile—the radial velocity of one peak coinciding almost precisely with

that of the Pleiades stars, and that of the

other peak being very nearly equal to the mean

radial velocity for the strongest components of

the interstellar lines. In all these investigations

of the fine structure of the interstellar medium,

one must insist on high angular resolution accompanied by narrow bandwidth; without

narrow bandwidth, the whole study of the

distribution of individual gas-cloud velocities

becomes impracticable. Once we have perfected our equipment, success in breaking down

the interstellar medium into its basic components should be within our grasp.

The greatest contribution so far to 21-cm

research by the 50-foot paraboloid of the Naval

Research Laboratory has been in the study of

absorption features observed in distant discrete

radio sources. Hagen, Haddock, McClain and

Lilley (Hagen, et al., 1955; Lilley, 1955) have

studied 21-cm profiles and related phenomena

in the continuum at wavelengths of 21 cm and

less for a number of discrete sources, some of

thermal origin, others not. The first relevant

21-cm profiles have exhibited certain amazing

absorption features which tend to confirm the

cloud or turbulent nature of the interstellar

gaseous medium. For the future study of

these profiles it will prove profitable to insist on

the highest attainable angular resolution accompanied by very narrow bandwidth of the

recording equipment. Almost every one of

these sources studied has presented us with a

different type of problem. The neatest absorption features are probably shown by the

Cassiopeia source, and it seems difficult to

escape the conclusion from the 21-cm observations by the Naval Research group and by

Davies (in press) and his associates at Jodrell

Bank that the source is in the Perseus Arm rather

than within 1,000 parsecs of the sun. The

study of the source at or near the galactic

center—or at least in the direction between our

sun and the galactic center—offers a host of

possibilities for future research and here very

much greater angular resolution than is attainable to date is urgently required. We shall

return in the next section to the Cygnus source,

which exhibits absorption features that originate far beyond the bounds of our own galaxy.

But the brief references that we have given

above indicate that 21-cm techniques offer

great opportunities for future research. The

contemplated interferometer for 21-cm research, with 2-component 90-foot steerable

paraboloids, should make possible an entirely

new approach to the study of discrete sources

and their absorption features. Large paraboloid antennas with apertures in excess of 100

feet also promise very significant contributions.

In cases where the 24-foot antenna is useless, the
NEW HORIZONS IN ASTRONOMY 145

50-foot antenna begins to suggest; but the 100-

foot antenna should be able to show the way and

the larger antennas should really be able to solve

some of the critical galactic astrophysical

problems of 21-cm research that are ahead of us.

Extragalactic research

To the Australian group at Sydney (Kerr,

Hindman, and Robinson, 1954) and their

associates (Kerr and de Vaucouleurs, 1956)

goes the credit for having detected and studied

the first 21-cm radiation from beyond our own

Milky Way system. Their studies of the

Magellanic Clouds, which have been blended

skillfully with the optical researches carried

on at Canberra (Buscombe, Gascoigne, and de

Vaucouleurs, 1954), have demonstrated conclusively that the 21-cm technique offers

entirely new possibilities for the study of

galaxies beyond our own Milky Way system.

These studies have already made important

contributions both to the dynamics of the

Magellanic Clouds and to our knowledge of

population distributions. Until it was found

that the 21-cm signals from the Small and from

the Large Magellanic Clouds are comparable

in strength, many had doubted that the Small

Cloud had anywhere nearly as much neutral

hydrogen as the Large Cloud. Through the

21-cm study we are learning a great deal about

the dynamics of the two Clouds. Their rotation seems well established, and the masses

derived from rotational velocities are in good

agreement with those estimated in other ways.

Studies of the distribution of random velocities

promise to add materially to our knowledge of

the dynamics of the gaseous interstellar medium

in these two galaxies.

The second big contribution to date to the

study of radiation from outside our galaxy

was made by Lilley and McClain (1955) in

their work on the principal absorption feature

in the discrete source in Cygnus. They have

established that an absorption feature in the

spectrum of the Cygnus source is present at

precisely the frequency predicted from the

optically known velocity of recession, produced

by the red-shift. The absorption feature

obviously originates in the galaxy inside which

the Cygnus source itself is imbedded. Not

only is the work of Lilley and McClain important in establishing the presence of a 21-cm

neutral hydrogen medium surrounding the

Cygnus discrete source, but the agreement

between the optically predicted red-shift velocity and the radial velocity determined by radio

techniques has important consequences for

cosmogonical theory.

The 21-cm radiation from the Andromeda

Nebula and other nearby galaxies remains as

yet undiscovered. Special equipment to aid

in its detection and measurement is now under

construction in Holland and elsewhere. The

importance of 21-cm studies for individual

galaxies of the Local Group can hardly be overrated. If we can attain angular resolution

sufficient to separate areas inside spiral arms

from those between the arms, very significant

contributions can be made to the solution of

the whole group of problems related to the

dynamics of galactic rotation and of spiral

structure. We shall then have also wholly

new evidence on the distribution of neutral

atomic hydrogen in spiral and other galaxies

and new light will be thrown on the problems

of stellar birth and evolution. For the future,

21-cm research of nearby galaxies, notably

those with spiral structure, has a high priority.

Quite recently Heeschen (in press) has succeeded in detecting 21-cm radiation associated

with a cluster of galaxies. At a frequency near

1387 mc/sec, corresponding to a recessional radial

velocity of close to 7000 km/sec, he found evidence for a weak but consistently present signal

(&T=2°K±0.°2) from the direction of the

Coma cluster of galaxies. The specifically

determined radial velocity of recession for the

same cluster is 6680 km/sec, which appears to

indicate satisfactory agreement. The dispersion in radial velocity derived from the 21-cm

observations appears to be of the order of 500

km/sec, somewhat smaller than the optical value

of about 800 km/sec. The total estimated mass

of neutral atomic hydrogen is a little less than

10" solar masses. Heeschen's result was obtained with the 24-foot radio telescope at

Agassiz Station. Information regarding the

distribution of neutral hydrogen inside the

cluster can be obtained only with instruments

of greater aperture, which afford higher angular

resolution.

A potentially very profitable approach to the
146 SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS

problems of intergalactic hydrogen is through

the study of absorption features. Field and

Heeschen have recently attempted to observe

in the spectrum of the Cygnus source absorption features produced by the general medium

of intergalactic hydrogen between the Cygnus

source and our sun. The results with present

equipment are inconclusive, but the analysis

of the available data shows that very rarefied

absorption features might be detectable, even

for the low densities that presumably prevail in

the intergalactic medium.

We should at this point refer again to the

interferometer array now under construction

by Bolton and his associates at the California

Institute of Technology. A variable-spacing

interferometer, with two basic components

of 90-foot aperture each, should make accessible for study a number of discrete sources

running literally into the hundreds. From

the absorption features, much information

should be obtainable—information relating not

only to the clouds between our sun and these

discrete sources but also concerning the distances of individual discrete sources.

Concluding remarks

It is now just about 5 years since Purcell and

Ewen first succeeded in detecting the 21-cm

line of neutral hydrogen. In the intervening

years much progress has been made, and I

doubt if many would have dared to predict in

1951 that by 1956 equipment equaling in excellence present-day radio telescopes and electronic

recording equipment would have become available at several institutions. But although we

may look with some satisfaction upon our

accomplishments to date, in reality 21-cm

research has only just begun.

I have purposely confined my remarks to

21-cm research in a most limited sense. I have

studiously avoided referring to the importance

of studies in the continuum at wavelengths of

50 cm and less, which must become an integral

part of all 21-cm studies. I have also omitted

all reference to other spectrum lines in the

microwave range, notably to the line of deuterium, the discovery of which would open up

an entirely new field for research of the interstellar medium by radio techniques. It would

be very helpful if, instead of having to confine

ourselves to the single 21-cm line, we could have

available additional lines, preferably those that

might give clues to the presence of atoms other

than hydrogen, and components of bands that

would permit the study of molecules, such as

OH and H2. The field of galactic and extragalactic microwave spectroscopy is indeed a

fruitful one for future exploration.

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SCHMIDT, M^ N O y

. Paper presented at Jodrell Bank Symposium 1955, in press. WESTEBHOUT, G.

SHKLOVSKT, I. S. • Paper presented at Jodrell Bank Sympo1949. Astron. Zhurn., vol. 26, p. 10. sium 1955, in press.

Some Observational Problems in

Galactic Structure

By Harold Weaver1

The past decade has brought many new insights into the problems of galactic structure.

New observational techniques—optical and

radio—have produced results that were thought

far beyond reach only a few years ago. We

now realize the importance of the various

structural units of the galaxy (clusters, associations, arms, etc.) as groups of evolving stars

of different ages. Data on the structure and

the kinematic properties of the galaxy provide

basic information for studies of stellar evolution; clusters and associations provide tests of

stellar-model evolutionary calculations. It is

clear that we must relate our galactic studies

to those of other nearby resolved galaxies.

Knowledge will advance most rapidly if we

understand how the structural details of our

galaxy relate to those of the sequence of galaxies.

New approaches

Several recent observations, to a considerable

degree, have reoriented the whole trend of

thinking about galactic structure and have

forcefully indicated the need for new approaches

to the problems. Among these observations

are the following:

1. Baade's (1944) decisive demonstration

that the space and physical-specification variables in the space-density function D(M, S;

R, 9, z) are strongly correlated.

The stellar space-density function D(M, S;

R, 0, z)dMdSdV will play an important part

in the discussion that follows; it indicates

symbolically the space density of stars in absolute magnitude range M to M-\-dM, spectraltype interval S to S-\-dS, and volume element

dV, centered at galactocentric position R, 0, z.

Other or additional variables can be used for

1

 Leuschner Observatory, University of California, Berkeley, Calif.

physical specification; color is frequently substituted for spectral type. If we fix the space

variables and represent the distribution graphically we obtain the Hess (1924) diagram.

The space distribution of nonstellar material

may be represented by Dt(R, 0, z), which refers

to particles of kind i.

For R small or z large in D(M, S; R, 6, z) one

finds pure or essentially pure Population II.

If to M, S are assigned ranges encompassing

intrinsically bright O and early B stars, Baade

found that in extragalactic systems the locus of

maximum of D(M, S; R, 0, z) traces out spiral

arms where Population I stars and much interstellar dust is to be found. Population I,

occurring in regions where z is small and R not

small, appears as an addition to, and intermingled with, the general substratum of Population II stars forming the disc of the system.

2. Morgan, Sharpless, and Osterbrock's

(1952a, 1952b) use of Baade's spiral tracer

technique to produce the first picture of spiral

structure in the galaxy.

3. Oort, van de Hulst, and Muller's (1952)

observation that the 21-cm radiation of hydrogen indicates hydrogen concentration in

spiral arms whose form and position confirm

and extend the work of Morgan and his

associates.

4. Ambarzumian's (see Kourganoff, 1951)

recognition of stellar associations and of their

importance in galactic and stellar evolution

problems.

In the following discussion we shall concentrate attention on the problems that fall

principally under 1 and 2; those that fall

principally under 3 and 4 will form the subject

matter of other papers. Effort will be made

to avoid duplicating material included in

149

381014—56 -11
150 SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS

Blaauw's (1955) excellent summary of the

International Astronomical Union symposium,

"Coordination in Galactic Research," which

should be read by all investigators in this field.

In this discussion we shall critically examine

the results of some recent researches in galactic

structure; and, rather than stress the agreements, we shall stress the discrepancies between

observation and theory and between the

observations themselves. While this approach

will accomplish our purpose—to expose areas

where problems need to be worked on—it

will show the field of galactic research in an

untrue light. The reader must guard against

gaining the false impression that little has

been accomplished in the field. The fact is,

rather, that much remains to be done.

Form of the arms

The form of the spiral arms first indicated by

Morgan and his associates, using emission

nebulosities, and by Oort, using hydrogen

concentrations, becomes increasingly complex

and less easily understood as additional observational data accumulate. Further work, both

observational and theoretical, in detailed delineation of the spiral arms by stars and by

gas is required. The Morgan-Meinel-Johnson

(1954) technique of very short spectra provides a powerful method of discovering distant

spiral tracers. Southern Hemisphere observations of the 21-cm radiation are especially

needed; 21-cm observations of high resolving

power in angle and in frequency are needed in

both hemispheres. Such observations will,

among other things, provide data for determining the forms of the gas condensations, and the

degree of association between OB stars and gas.

The distance scale within the galaxy

The distance scale presently employed in tracing

spiral structure by means of OB stars is correct

only to the extent that (1) absolute magnitudes

and colors of the MK classification system are

correct, and (2) the ratio of total interstellar

extinction to color excess is known.

The distances of the hydrogen condensations

are correct only to the extent that (1) an

accurate differential-rotation function Aw(R) is

available, and (2) the interstellar gas is free

from large-scale peculiar motions.

Stellar distance scale.—The present scale of

absolute magnitudes (Keenan and Morgan,

1951) is not entirely independent of galactic

rotation constants (first order theory only).

These, in turn, are dependent upon an older

absolute magnitude scale and on corrections

for average interstellar extinction. Work of a

fundamental nature is here necessary to eliminate the possibility of a closed circuit of

systematic errors that do not easily reveal

themselves. Fundamental absolute-magnitude

scales are required—scales derived from stellar

associations, clusters, moving clusters, and

stars of large parallax (with the latter two

types of objects serving for establishment of

zero point). The degree of agreement between

main sequence magnitudes (and also luminosity

class IV, III, . . ., magnitudes) derived from

different associations, clusters, and so forth,

will be of extreme interest; comparisons with

field stars in the solar neighborhood will likewise be of importance. Several investigators

are engaged in observations of this type and

are producing very interesting results.

Observations of the ratio of total interstellar

extinction to color excess are contradictory.

On the basis of both photographic and photoelectric studies (Baade and Minkowski, 1937;

Stebbins and Whitford, 1945; Sharpless, 1952),

the Orion Nebula has been pointed out as

possessing an extinction law appreciably different from that of other interstellar material.

A photographic spectrophotometric study of

O and B stars by Divan (1954) indicates

precisely the same extinction law in the Orion

region as in all other regions investigated.

On the other hand, very accurate photoelectric

observations by Johnson and Morgan (1955)

indicate an "abnormal" ratio of extinction to

color excess for the Cygnus region in the sense

that relatively more ultraviolet extinction is

present in that region than in the CepheusMonoceros region. Space variations of the

extinction law complicate photometric distance calculations and pose difficult problems

for the theoretician attempting to explain the

distribution of particle size in the interstellar

material. The observational determination

and subsequent theoretical discussion of the

wavelength dependence of interstellar extinction requires further attention.
NEW HORIZONS IN ASTRONOMY 151

Hydrogen distance scale.—The differential

rotation function Au(R) for R^>R0 must be

derived from stellar motions or from the mass

distribution in a model galaxy; circular motion

is assumed in either case. If derived from

stars, Aco(i?) is dependent upon the accuracy

of the stellar distance scale. For R<C.R0 the

hydrogen observations themselves permit, in

principle, the establishment of Aw(R).

A 21-cm observation in the direction of the

inner part of the galaxy yields a »max which, in

a galaxy in pure rotation and with a monotonic

w(R) function, arises from the point that, on

the given line of sight, is nearest the galactic

center. Precisely,

(1)

lia Sin A

In practice, there may be observational

difficulties in the measurement of the »„,„

described in equation (1) because of local

peculiar motions of the gas, random motions

or cosmic dispersion within the gas, and

variations of gas density along the line of

sight.

For R—Ro small, and with the definition

A=— )£ Rou'o, equation (1) may be approximated by the expression,

Ro=- 2A(1 —sin X) sin X (2)

On the basis of linear theory one can find

Ro, provided vmn has been measured and the

Oort A parameter is known.

Van de Hulst, Muller, and Oort (1954) used

equation (2) to compute Ro with an adopted

value of .4=19.3 km/sec per kiloparsec, which

is a mean of various A values derived on linear

theory. It is likely that A is overestimated.

(Weaver, 1955a, 1955b, 1955c.) From four

»max values they found i?0=8.26 kpc; they employed the value i?0=8.2 in further calculations

of Aw(i?) from equation (1). The scale of

is fixed by Ro; systematically, the

or <a(R) curve found from the »„„«

observations is, then, no more precise than the

Ro or, if we look back one step in the reasoning,

than the A value and the linear theory on

which it was based.

Van de Hulst, Muller, and Oort checked Ro

against Baade's (1954) value for the same

881014—58 12

parameter, 8.16 kpc, derived from the distribution of RK Lyrae variables in the region of the

galactic nucleus. The excellent agreement of

the two values does not necessarily confirm

their correctness. Baade's estimate could be

appreciably changed by a small revision of the

interstellar extinction correction employed. A

change of 0.25 mag in the extinction would alter

Ro by a kiloparsec.

The crucial need of an accurate Ro is apparent; further investigation is imperative.

The uncertainty in the galactic distance scale

is also indicated by inconsistency in the currently quoted values of Ro, «o, and Fo. Van de

Hulst, Muller, and Oort (1954) found «o=26.4

±2.5 (m. e.) km/sec per kiloparsec, a value

derived by differencing A and B. Weaver and

Morgan (1956) employed proper motions of 79

Cepheids in mathematically unbiased equations

to find «o=23.2 ±5.0 (m. e.) km/sec per

kiloparsec. These values combined with Ro=

8.2 kpc yield, respectively, Vo=216 km/sec

and 190 km/sec. These are not fundamental

determinations since they represent the product

of two independently derived quantities.

Mayall (1946) found Fo=200 km/sec from

radial velocities of 50 globular clusters. This

is a fundamental determination, though it

probably does not represent the true value of

Vo, being smaller than that value by the

difference between the circular velocity of the

galaxy at Ro and the corresponding value for the

system of globular clusters. A direct observational determination of this difference does not

exist. Mayall (1946) attempted to evaluate

it by determining the solar motion with respect

to the extragalactic nebulae. From several

solutions he concluded Vo=300 km/sec. Humason and Wahlquist (1955) have derived an

essentially identical value. Any peculiar motion of the galaxy is included in this result,

which therefore does not permit a strong conclusion that Vo is as great as 300 km/sec, but

strongly indicates that F0>200 km/sec. Fricke's

(1949a, 1949b) value of Vo from high velocity

stars, also a fundamental determination, likewise indicates F0>200 km/sec. Drawing an

analogy between our galaxy and the Andromeda

Nebula, we reach a similar conclusion from

Mayall's (1951) velocity measurements in

Andromeda. An inconsistency thus appears.
152 SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS

The product J?o«o is smaller than the Vo value

derived by independent fundamental methods.

Examination of Ro and w0 is required. Of the

two quantities, Ro is, percentagewise, the more

uncertain.

The u(R) curve derived by Kwee, Muller, and

Westerhout (1954) from 21-cm observations

shows waves which, they suggest, represent

departures from circular motion which are

correlated with position with respect to spiral

structure. The general question of systematic

deviations from circular motion over sizable

GALACTIC

CENTER

gen. For the stars, accurate photometric distances and radial velocities can be derived for

objects nearer than r0, an observationally set

limit. Over the region of space thus defined

(shown in fig. l,a) the motion pattern can, at

least in principle, be derived and regional departures from circular motion deduced. Hydrogen observations provide no distance scale

independent of velocity; they permit study of

departures from circularity only over the volume of the galaxy in which tWx occurs (shown

in fig. 1,6) in the range Ri<R<Ro. The

\ CENTER /

FIGURE 1.—The plane of the paper represents the plane of the galaxy, a, Crosshatched area indicates the region over which stars

permit the study of departures from circular motion, b, Crosshatched area indicates the region over which 21-cm radiation

of hydrogen permits study of departures from circular motion.

volumes of the galaxy merits investigation on a

broad basis. Observations of both hydrogen

and stars are desirable. Of particular interest

will be tests for similarity of kinetic properties

of stars and gas in the same volume of space and

for a correlation between gas and stellar motions

and spiral structure as suggested by Kwee,

Muller, and Westerhout.

Study of departures from circular motion is

likely to prove easier with stars than with hydrogalactocentric distance Rx is fixed at the point

at which strong turbulence in the interstellar

hydrogen sets in. In the case of hydrogen, a

large peculiar motion occurring in the line of

sight, but not within the indicated volume,

could obscure a smaller peculiar velocity within

the indicated volume and thus confuse the observation. To a lesser degree, variations in the

distribution of random velocities in the hydrogen also obscure the situation. Such difficulties
NEW HORIZONS IN ASTRONOMY 153

do not exist for the stellar observations. Hydrogen observations, on the other hand, extend to

a much greater distance than do stellar ones.

Comparison of flmax values derived at corresponding angles on either side of the galactic

center will be of very great interest.

Association of spiral tracers and hydrogen

No adequate test has yet been made of the commonly accepted hypothesis that arms or structures delineated by very early type stars are

coincident in space with hydrogen gas density

maxima. There are indications that these may

not always be coincident (Weaver, 1953; van de

Hulst, Muller, and Oort, 1954). OB stars may

appear where no hydrogen density maximum

occurs.

Several tests of this hypothesis can be proposed. The simplest involves investigation of

the simultaneous occurrence of hydrogen maxima and concentrations of early type stars over

the surface of the sky. This test was independently proposed by W. W. Morgan during

a colloquium given at Berkeley, Calif., in November 1955. An alternative test involves the

agreement of velocity maxima of the radial

velocity distributions of the OB stars (treated

as particles in a star gas) and of the hydrogen

gas, properly weighted to match the distribution of the stars in longitude. A third test

involves the determination of the closeness of

association in space of maxima in the space

density functions of OB stars and of interstellar

hydrogen as found from the 21-cm observations. Distances of the stars are to be derived

photometrically; distances of the hydrogen

maxima are to be determined from a rotation

function based on the same photometric distance scale used for the stars. Different

assumptions are involved in the three cases;

all three tests should eventually be made.

Composition of the arms

Understanding of the physics of the spiral

arms will require, among other data, knowledge

of their composition. Nearby extragalactic systems will undoubtedly play an important role

in composition studies. Photographs of M 51

published by Zwicky (1955) are instructive;

analogous photoelectric measurements would be

invaluable. Isophotes and isocolor contours,

even of a rudimentary kind, for details in a few

properly chosen extragalactic systems would be

of great interest. A start on this problem

could be made by measuring magnitude and

color differences between arm and interarm

regions of some of the most resolved extragalactic systems.

Spiral tracers are highly concentrated in the

arms; data on the space distribution of latertype Population I stars, B 3, B 5, . . . , in and

near the arms is lacking. We have no knowledge of how prominent the arms would appear

if the galaxy were viewed through a series of

hypothetical filters, each of which would permit

us to see stars in only a small spectral-type

range, B 3, B 5,. . . . The ordinary method of

analysis of star counts does not offer much

promise of producing such a picture of the

galaxy. Rather, we should obtain the picture

with the aid of accurate individual photometric

distances derived for each member of a large

sample population of stars. An empirical model

of the space distribution of any one particular

type of star could then be constructed. Data

such as density gradient perpendicular to the

galactic plane, in the galactic plane, and perpendicular to the arm, and so forth, could be

inferred from the model for each stellar type.

An approach through such an empirical

model avoids the concepts of a continuous

space-density law and of a space-invariant

luminosity function. It does not involve the

solution of an integral equation for the spacedensity law or the correction of the apparent

space-density law with the aid of an average

interstellar extinction relation in a specified

direction. Each star is treated individually;

maximum resolution is achieved. A limit of

completeness in absolute magnitude M(r) can

be specified at each distance; the space density

of intrinsically faint stars is not inferred from

a few intrinsically bright ones as when the

integral equation is solved.

Extensive observations are required for such

a program; a cooperative venture involving

several observers and telescopes might prove

advantageous. Stromgren's (1951, 1954) accurate photoelectric method of spectral classification should play an important part in any

such plans. Stromgren's results indicate that

the photometric probable error of one observa-
154 SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS

tion of a predicted absolute visual magnitude

is less than 0.2 mag; in distance estimate a

probable error of approximately 15 percent

might be expected. With the 82-inch

McDonald telescope, observations could be extended to apparent magnitude 14. Further

exploratory work in the use of this important

observing technique is called for. The results

obtained should provide data for the establishment of an empirical space distribution as a

test.

The need of an accurate and fundamental

calibration of the absolute magnitudes and

colors of the MK spectral types and luminosity

classes as well as an accurate value of the ratio

of total extinction to color excess is reemphasized

by a program of this type.

Variations in the arms

Within the spiral arms are found variations in

the space density of stars and in the stellar

mixture. Observational data leading to a

better understanding of such variations would

be of considerable value.

A pronounced example of the first kind of

variation is found in the longitude range 110°

to 145° adjoining the rich Perseus star cloud.

In this range of longitude there exists a deep

minimum in the number of OB stars (Weaver,

1953); there is essentially a discontinuity in

the Perseus arm as delineated by early-type

stars.

Variations in the content of the arms frequently appear. The Perseus cloud contains O

stars and late-type supergiants (Bidelman, 1943,

1947); classical Cepheids are found within it.

The Orion association contains principally main

sequence stars of type O 6 and later; it contains

no supergiants of type later than B (Sharpless,

1952, 1954). The Scutum association, which

has M 11 as its nucleus, contains main sequence

stars of type B 9 and later, and ordinary giants

(Russell, 1953). Gurzadian (1949) has pointed

out many cases of this sort.

A preliminary study of variations in the

arms might be made on the basis of galactic

clusters and Trumpler's cluster types.

Luminosity functions and related quantities

An observing program of the type outlined in

the preceding sections would provide detailed

information on the space variation of D(M,

S) R, 6, z) over the M-range and r-range

accessible to observation. However, such a

program now represents only a possible future

development; information on stellar populations which will provide data for stellar evolution studies must be provided more rapidly.

Among observable quantities that provide

such information with less completeness than

D(M, S; R, 6, z) but that are much more

easily obtained are the following:

1. The luminosity function, L(M; R, 6, z),

which is the integral of D with respect to S.

2. Integrated magnitudes, colors, or spectra,

which represent, essentially, integrals of the

luminosity function over specific wavelength

ranges and which refer to a defined volume

element of the galaxy or of some extragalactic

system.

3. The ratio of number of stars brighter than

some fixed M and within a defined volume

element to the total luminosity of all stars

within the same volume element (Baum and

Schwarzschild, 1955). The wavelength ranges

in which M and the total luminosity are measured must be specified.

4. The ratio of mass to luminosity (measured

in a specified wavelength range) of the material

within a defined volume element of the galaxy

or of an extragalactic system.

5. The ratio of luminosities of stars contained within some defined volume clement of

the galaxy or an extragalactic system measured

in several different specified wavelength ranges

to the luminosity of Uio same stars in some

standard wavelength range.

Basically, the integrated quantities represent

certain average properties of the stellar population in the defined volume element. The ratios

compare one portion of the population with

another portion of the same population. The

particular ratio taken fixes the portions of the

populations compared. All of these observable

quantities need further investigation and application within our own galaxy and extragalactic

systems.

The luminosity function of stars characterizing the solar neighborhood has been determined

by van Rhijn (1936), by Luyten (1939, 1941),

and others. Van Rhijn's function shows a

minimum at absolute magnitude -f 8; Luyten's
NEW HORIZONS IN ASTRONOMY 155

does not. The reality of this feature requires

checking. The test might be made with the

aid of late-type dwarfs discovered spectroscopically (Vyssotsky, 1943, and coworkers,

1946, 1952). With the present functioning

of Schmidt telescopes, intrinsically faint stars

discovered spectroscopically should play an

increasingly important role in the determination of L(M; Ro, 0, 0).

The van Rhijn luminosity function presumably characterizes Population I; the exactness of the characterization may bear

further analysis. The method of deriving

the luminosity function entails the use of

volume elements of different sizes for different

luminosity ranges. The very bright stars are

observed over a volume element of great size,

a cylinder of possibly a kiloparsec radius; the

very faint stars are observed within a spherical

volume element of a few parsecs radius centered

on the sun. The highly luminous stars are

known to be nonrandomly distributed in the

galactic plane; indications are (McCuskey,

1951) that the fainter stars also vary in mixture

over moderate distances from the sun. The

van Rhijn luminosity function was derived when

there was still belief in the universality of the

function. There is no certainty that it refers

to a homogeneous population of stars. The

function is, for the faint stars, essentially

L(M; R, 0, 0); for the bright stars it is some

average L(M), the average being taken over

a large and somewhat uncertainly specified

volume element.

An important step towards a clearer understanding of the Population I luminosity function could be taken by determining the luminosity functions of a number of representative

Population I objects; namely, galactic clusters,

star clouds, associations. Such luminosity

functions could be determined with adequate

accuracy by statistical means, provided an

accurate magnitude sequence was available for

each case. Of particular interest would be a

study of variations in the form of the luminosity

functions for stellar groups of different Trumpler type or different age, different size, and so

forth. The possibility of forming from such

a sample a mean function representative of

Population I in the solar vicinity should be

investigated. Variations in different groups

or changes through evolution may, of course,

make the establishment of such a mean function

impossible on the basis of present knowledge.

Luminosity functions of typical Population

II would also be of great interest. Important

results for luminosity functions of globular

clusters have recently been given by Sandage

(1954) and by Tayler (1955); for nebulae by

Baum and Schwarzschild (1955), by Baum

(1955), and by Roberts (1955, and in press).

For these latter cases indirect methods using

integrated magnitudes and multicolor observations were employed.

Variations have been found in luminosity

functions of typical Population II systems just

as they have in Population I systems. An

elliptical nebula appears to contain a much

larger percentage of dwarfs than a globular

cluster. The luminosity function of the dwarf

branch of a globular cluster appears remarkably

similar to that of the stars in the solar neighborhood; a typical galactic cluster contains a

much smaller percentage of dwarfs than does

the solar neighborhood. Comparisons of this

sort based upon many more cases than are now

available will be of very great interest.

They may well provide the key to an understanding of the mixture of stars we find in the neighborhood of the sun, particularly if information

of the space density of stars of different spectral

types and luminosity classes in and between

the arms can be derived.

The local system

Interpretation of space density results in the

region of the sun in terms of a simple picture

of spiral arms may prove difficult because of

the interference of the "local system" (see Bok,

1937, for earlier references). The observations

by Nassau and Morgan (1951) demonstrate

that the distribution of bright OB stars follows

the outline of the local system. Heeschen and

Lilley (1954) have found the local system in

the 21-cm radiation of hydrogen. The structure likewise appears among both the bright and

dark nebulae. A thorough modern study of the

local system would be of interest; it would also

be of value as an aid in disentangling the

local system and what one ordinarily thinks of

as spiral structure.
156 SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS

Other aspects of the galaxy

For an understanding of the galaxy a knowledge of the interstellar material is no less

important than that of the stars. Problems

of some aspects of interstellar matter have been

described where necessary: the degree of association of interstellar hydrogen and spiral

tracers, motions of the hydrogen, and the value

of the ratio of interstellar extinction to reddening. However, general problems of the

interstellar medium and of the 21-cm radiation

of hydrogen fall outside the scope of this

review. They will be found in the reviews of

those subjects. Likewise, problems of motions

within the galaxy have been introduced only

where necessary since they form the subject

matter of a separate paper. In actual practice,

the study of the distribution of stellar or

interstellar material and the study of kinematics

cannot be separated if a complete view of the

problem is to be gained. The distribution of

material and the motions of the material are

component parts of one problem that must be

considered as a whole.

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1951. Pub. Obs. Univ. Michigan, vol. 10, p. 43.

OORT, J. H., VAN DE HULST, H. C, AND MuLLEB, C. A.

1952. Trans. Int. Astron. Union, vol. 8, p. 512.

ROBERTS, M. S.

1955. Publ. Astron. Soc. Pacific, vol. 67, p. 327.

. Astron. Journ., in press.

RUSSELL, J. A.

1953. Astron. Journ., vol. 58, p. 89.

SANDAGE, A. R.

1954. Astron. Journ., vol. 59, p. 162.

SHARPLESS, S. L.

1952. Astrophys. Journ., vol. 116, p. 251.

1954. Astrophys. Journ., vol. 119, p. 200.

STEBBINS, J., AND WHITFORD, A. E.

1945. Astrophys. Journ., vol. 102, p. 318.

STROMGREN, B.

1951. Astron. Journ., vol. 56, p. 142.

1954. Astron. Journ., vol. 59, p. 193.
NEW HORIZONS IN ASTRONOMY 157

TATLEB, R. J.

1955. Astron. Journ., vol. 59, p. 413.

VAN DB HULST, H. C , MuLLEH, C. A., AND OOBT, J. H.

1954. Bull. Astron. Inst. Netherlands, vol. 12,

p. 117.

VAN RHIJN, P. J.

1936. Publ. Astron. Lab. Groningen, No. 47.

VTSSOTSKT, A. N.

1943. Astrophys. Journ., vol. 97, p. 381.

VTSSOTSKT, A. N., JANSSEN, E. M., BT AL.

1946. Astrophys. Journ., vol. 104, p. 234.

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1952. Astrophys. Journ., vol. 116, p. 117.

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1953. Astron. Journ., vol. 58, p. 177.

1955a. Astron. Journ., vol. 60, p. 202.

1955b. Astron. Journ., vol. 60, p. 208.

1955c. Astron. Journ., vol. 60, p. 210.

WEAVER, H. F., AND MORGAN, H. R.

1956. Astron. Journ., vol. 61, p. 268.

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1955. Publ. Astron. Soc. Pacific, vol. 67, p. 232.

Galactic Clusters and Associations

By A. Blaauw

 l

Research on galactic clusters and associations

can be divided into two categories: the properties of the individual objects, and the properties

of the system of the objects as a whole. The

first aim of these investigations is, of course, to

arrive at a knowledge of the present condition

of the individual clusters and associations and

of their present motions and distribution in

space. But the ultimate purpose is to understand these conditions as a result of the origin

and evolution of the clusters and associations.

This problem is immediately connected with

that of the history of the galaxy and with that

of stellar evolution. This double aspect makes

the study of clusters and associations particularly significant.

There is no sharp division between galactic

clusters and associations. This is best illustrated by the fact that certain objects, such as

NGC 2244, may be classified under either of

these categories. Most associations must have

short lifetimes and will have dispersed among

the general population of the galaxy within one

or two revolutions of the galaxy (a few hundred

million years). Galactic clusters are more

stable objects but this need not mean that their

origin is essentially different from that of the

associations. They may simply be remnants

of the rich central parts of some associations with

very small internal motions, that had a negative

energy from the outset, or have disposed of a

small positive energy by the ejection of part

of their stellar content. Stellar associations

offer an opportunity to study the properties of

the very youngest stars, and they seem to be

one of the best places to look for information

on the process of star formation itself. Galactic

clusters offer a variety of star samples, each of

which informs us about the properties of stars

formed simultaneously and at the same location in the interstellar medium, under certain

»Yerk«s Observatory, Williams Bay, Wls.

common circumstances. A vast amount of

observational data has been collected in the

past for the galactic clusters, most of which will

serve as a basis of more refined work made

possible by the improved present and future

observational techniques. Much less has been

done on the associations, whose importance

has only recently become clear. However,

what has been done observationally so far,

especially in the field of photometry, meets the

standards required for our present demands.

Having in mind that research on galactic

clusters and associations can contribute to our

understanding of the evolution of the galaxy

and of the stars, we shall consider three major

trends in this research, in the three sections

following.

Individual properties of the stars

Individual properties of the stars in clusters

and associations may be: the absolute magnitude, the color, the spectral type, the angular

momentum, the duplicity or multiplicity, and

the motion of the star with respect to the center

of the cluster or association. Such studies are

important because the stars in the cluster or

association can be assumed to be of approximately equal ages, to have formed at the same

location in the interstellar medium, and,

perhaps, to have uniform chemical composition.

This last assumption is uncertain since the

abundance of interstellar grains may have

varied within the cloud complex in which the

stars were formed. In principle, knowledge of

these properties of the stars in a large number

of clusters and associations would provide an

ideal basis for studying (1) the properties of

stars of equal age but of different mass and

angular momentum; (2) the changes, with age,

not only of the stars' position in the HR diagram but also of their rotational and duplicity

properties.

159

381014—56 -13
160 SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS VOL.1

Trumpler's extensive investigations of galactic clusters have revealed the large variety of

patterns in the HR diagrams and this, combined

with modern concepts of stellar evolution,

shows that these objects represent a wide range

of evolutionary stages. The modern equivalent

of these diagrams, based on accurate photoelectrically measured colors and magnitudes,

has already allowed more detailed discrimination of the properties of different stars in the

same cluster, and more refined comparison of

the color magnitude arrays of different clusters.

It has, moreover, the great advantage of reaching considerably fainter apparent—and, hence,

also absolute—magnitudes than the spectroscopic method.

While color and absolute magnitude are the

most obvious quantities to choose in defining

the physical properties of a star, they do not

describe them completely. Initial chemical composition and angular momentum cause variations of the array in the color absolute-magnitude diagram for clusters of equal age. We

should consider also the duplicity or multiplicity

of the stars and their motions in the cluster.

These will probably have no relation to the

color-magnitude arrays, but they must have an

important bearing on the process of the formation of the cluster from the interstellar medium

and will have to be taken into account in subsequent studies of the evolution of the clusters.

There may be other properties of the stars in addition to those listed above, which, though not

yet easily accessible, may provide important

clues; the magnetic fields might be among these.

There can be little doubt that the first desideratum for future work on galactic clusters is the

determination of accurate color absolute-magnitude arrays for a large number of clusters.

These should provide the basis for subsequent

studies of the secondary effects mentioned

above, which may be feasible only for a limited

number of stars and a limited number of clusters. These latter could then be selected so as

to represent the most useful variety of evolutionary stages.

A critical step in the determination of the

color absolute-magnitude arrays is the conversion of apparent magnitudes into absolute magnitudes. There are only a few clusters for which

we have geometrically determined distances:

the Hyades, Ursa Major, and a Persei clusters.

The use of differential galactic rotation as an

indicator of distance has not produced valuable

distances in the past, but it may do so in the

future for the young clusters if these follow the

spiral structure of interstellar hydrogen closely

and if radial velocities determine uniquely to

which arm the clusters belong. H. L. Johnson's

approach to the problem of the absolute magnitude determination, based on the assumption of

identical "zero-age main sequences," may prove

to be very useful provided sufficient evidence

can be produced to show that the chemical

composition does not affect the position of the

main sequence. One must bear in mind, however, that this approach eliminates the detection of composition differences unless these are

of a more pronounced character in the rapidly

evolving part of the color absolute-magnitude

diagram than in the slowly evolving part; the

determination of the absolute magnitude scale

should be based, preferably, on the latter.

The extension of the color-magnitude arrays

to faint stars may require the identification of

fainter cluster members than are known at

present. This holds even for the nearest

clusters, the Hyades being a good example of

a cluster whose faint members are still unknown.

It will in general be the case for the poorer

clusters observed against the rich background

of the Milky Way. So far, observations of

the faint stars in clusters have usually been

confined to a sample considered representative

as far as the magnitudes and colors are concerned, with little attention being given to the

estimation of the total numbers of faint members. Yet, this is an important item if we

want to consider the density of the population

along the main sequence at zero age for different clusters. That differences in this respect

exist follows from the comparison of NGC

752 and Praesepe. Such differences are probably related to the mass density and the state

of motion in the interstellar medium from which

the cluster originated. Their observation may

provide interesting data on the formation

process of the clusters.

Existing data on the identification of faint

cluster members can in many cases be considerably increased by the measurement of new

proper motions in the fields of the clusters.
NEW HORIZONS IN ASTRONOMY 161

The large collections of photographs of galactic

clusters, taken at several observatories 30 or

more years ago, are becoming more and more

useful for this purpose. In many cases, work

of this kind done in the years preceding World

War II cannot be considered to be up to date

for the future projects mentioned above.

Comparison of the old epoch plates with modern

ones may give proper motions of several to

10 times the weight of those available at present.

It is gratifying to know that several observatories that possess fine collections of old plates

have already started comparing these with

modern ones, whereas others have expressed

their interest in undertaking programs of this

kind. International Astronomical Union Commission 37 on star clusters and associations at

the Dublin Assembly has adopted a resolution

to the effect that observatories which possess

collections of early epoch cluster plates present

lists of these to the president of Commission

37 and indicate to what extent they are in a

position to collaborate on proper motion work.

Observations of the stellar rotation velocities

and of the frequency of spectroscopic and

visual binaries have, in the past, been given

relatively little attention. Yet they would seem

to offer most interesting prospects for future

research. Why are there few or no spectroscopic binaries among the Pleiades although

visual doubles are plentiful, whereas spectroscopic binaries are frequent in other clusters,

notably in those containing many early-type

stars? Are we facing here an evolutionary

effect or is this difference directly related to

different initial conditions in the interstellar

medium at the time the clusters were born?

Whatever the answer may be, it is obvious

that a systematic study of the frequency of

rotational velocities and of the duplicity

characteristics in clusters of a variety of ages

and appearances will ultimately give valuable

information on the formation problem.

The foregoing remarks referred in the first

place to the work on galactic clusters. Although

the basic problems are very much the same for

the associations, the emphasis in future work

will probably be somewhat different. The

associations on the whole are much younger

than the clusters so that the stars have evolved

very little, while the internal velocities are

generally large and very likely much the same

as at the time of formation of the stars; therefore, they can indicate the initial properties

of the stars and may yield clues to our understanding of their process of formation. The

observations made so far have been concerned

mainly with the photometry (absolute magnitudes and colors), the proper motions, and

radial velocities, while a beginning has been

made with the systematic investigation of the

duplicity properties. But all these referred

only to the brightest part of the population and

the stars of spectral types about B5 and earlier.

Very little is known as yet of the occurrence

of fainter, later-type stars. Stars down to

spectral type A0 along the main sequence

have been surveyed in the Orion association, in

III Cep, and in the most concentrated parts

of a few other objects. But in two of the

nearest associations, II (f) Perseus and I

Lacerta, we are still ignorant about the members of types B5 and later. Yet these two

associations are favorably located at relatively

high latitudes where the confusion by the

general galactic background is less serious than

in most other cases. Surveys of the late B and

later-type stars in associations will suffer from

a few drawbacks: the area over which the stars

have to be classified is much larger than for

the clusters—for the nearest associations of

the order of 100 square degrees—and, since

the internal velocities are large, members and

field stars cannot be distinguished uniquely.

The membership may have to be based on

statistical evidence, like that, for instance,

showing that the concentration in the sky for

a particular type of star is greater in the region

of the association than outside it. The size

and direction of the proper motions may help

when the expanding character is pronounced,

but then it can be used only in the outer regions

where the density of the population is small.

In spite of these difficulties, knowledge of the

fainter membership of the associations must

be considered one of the most urgent goals of

research on these objects. The first approach

should, of course, be by Schmidt telescopes or

large-field astrographs equipped with objective

prisms.
162 SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS

Dynamical studies

Dynamical studies of clusters have, in the past,

been mainly concerned with systems in a

state of approximate equilibrium, and have

dealt with the relation between internal velocity

distribution and mass distribution, the times

of relaxation, and the rate of escape of stars

from the cluster. Investigators usually assumed that the clusters are approximately the

age of the universe, several billion years.

Recently acquired knowledge of the ages of

clusters derived from increasing information

on stellar evolution offers the possibility for

a fresh approach to dynamical studies of

clusters. Various questions come to our mind.

We notice that the distribution of stars in the

Hyades is quite irregular. We estimate its

age at about 109 years. Can we reconcile

these two observations, estimate how many

stars must have escaped from the Hyades, and

predict its future constitution? Suppose we

indeed collect data on the ages of a large number of clusters as suggested in the preceding

section, can we explain the numbers of clusters

of various ages as a result of the predicted rate

of disintegration ?

The dynamics of associations presents problems of a somewhat different kind. Especially

in the youngest objects, the forces acting on the

stars are probably mainly due to the interstellar

matter present in the region of the association.

If we make the reasonable assumption that

only a small fraction of the interstellar matter

has condensed into stars, its mass may very

likely exceed the total mass of the stars many

times. Probably the most promising aspect

of the study of the motions in associations is

that it will give us direct information on the

state of motion of the generating interstellar

cloud complex, and it is intimately connected

with the complicated problems of the dynamics

of cosmic clouds. It is clear, however, that we

must obtain the most precise data on the

motions of individual stars. For the brightest

ones (m<C6) we will depend on continued

meridian observations. For the fainter stars

(around m=9) we must rely on photographic

determinations of new positions. These will

require special programs involving large numbers of plates covering the limited regions of

the associations in order to attain high accuracy

for the individual stars. In addition to the

proper motion work we must consider radial

velocity observations as one of the principal

desiderata for future work. This has been

mentioned also hi connection with the detection

of spectroscopic binaries.

The system in the galaxy

Investigations of the space distribution and the

kinematics! properties of clusters have shown

that they occupy a flat subsystem in the galaxy

with peculiar motions of the order of 10 km/sec.

Only in recent years have the more accurate

photometric distance determinations revealed

the connection between open clusters and the

spiral structure.

For future work in this field, also, the accumulation of a large number of accurate colormagnitude diagrams with the inferred knowledge of the ages of many clusters forms the

basic requirement. Once this information has

been obtained one can study the distributional

and kinematical properties for different age

groups. For the youngest clusters the evidence

already obtained for the distribution along the

spiral arms will be strengthened, and it will be

interesting to see where, along the spiral structure of interstellar gas, the formation of clusters

used to take place. Data on proper motions,

already referred to, and mean radial velocities

will provide information on the velocities of

these objects at the time of formation.

The most intriguing of these studies may,

however, well be those dealing with the somewhat older objects. Shall we be able to reconstruct from their present distribution and

motions their distribution at the time of formation ? This may be possible for the slow moving

clusters whose ages are now of the order of one

revolution of the galaxy (200 million years).

If we assume their space velocities to be of the

order of 5 km/sec, they have traveled during

their lifetime a distance of about 1,000 parsecs

from their origin. This is probably not too far

to allow an approximate reconstruction of the

distribution of the original locations. Thus we

may be able to study the secular changes in the

spiral structure of the galaxy, a problem which

we otherwise could tackle only by studying the

space distribution of the A-type stars. These
NEW HORIZONS IN ASTRONOMY 163

are, of course, much more numerous, but the

clusters have the advantage of being detectable

over a larger area.

These studies require measurements of the

proper motions and of the radial velocities of

the brightest stars in the clusters, in addition

to the photometric data. It may not be too

optimistic to expect that present and future

observational equipment may collect these data

for several hundred clusters.

Studies of the space distribution of the associations are well under way at present and are

of great interest because they outline the spiral

structure. Also radial velocities of the association are being measured and they are essential

in supplementing the studies of spiral structure

based on the 21-cm hydrogen emission. Work

on the proper motions, of special value for the

nearer objects, will be greatly facilitated by the

recently planned repetition of the AGK2.

Globular Clusters Observed through

a Crystal Ball

By W. A. Baum1

At the focus of a large telescope, a globular

cluster is one of the most magnificent sights in

the sky. One can see thousands upon thousands of stars all crammed together within a

few minutes of arc. Embedded like jewels here

and there in the sparkling swarm, the brighter

stars run the full gamut of color; the brightest

are brick red while others, slightly less bright,

are distinctly blue. Toward the center of the

cluster, the population becomes fantastically

dense. It is amusing to consider what our sky

would look like if the earth were associated

with a star in the nucleus of a globular cluster.

Perhaps the best way to judge what the

future has in store will be to examine the

present state of our knowledge and some

present lines of speculation to see the gaps

which need to be filled.

Globular clusters constitute a very well defined class of objects, all pretty much alike in

their general features. You could hardly

mistake one for any other type of system. Out

of more than a hundred observable globular

clusters in our own galaxy, only two or three are

still debatable cases.

The distribution of globular clusters in the

galaxy is almost spherically symmetric about

the galactic nucleus, whereas stars like those

in the solar neighborhood, along with gas and

dust, are restricted to a highly flattened diskshaped region. This difference in distribution

is associated, as one might expect, with a

difference in angular momentum; globular

clusters probably do not as a group share to any

large degree in the rotation of the galactic disk,

but they have large individual radial velocities.

Except in the immediate region of the galactic

nucleus, the number of globular clusters per unit

1

 Mount Wilson and Palomar Observatories, Pasadena, Calif.

volume decreases as the inverse third or fourth

power of the distance from the nucleus, and

there is no well-defined outer boundary to their

domain. It would be very valuable to have a

more complete and accurate analysis of the

distances and velocities of all known globular

clusters in the galaxy, but the task is a big one.

The stars in globular clusters are an entirely

different breed from those with which we are

familiar in the solar neighborhood. They have

different chemical abundances and they consequently follow different patterns of evolution

as they age. These differences can be deduced

from spectra and from color-magnitude diagrams. Stars of the globular-cluster kind are

sometimes referred to as Population II. They

are found not only in the clusters but also as

unattached vagabonds wandering through the

same regions of the galaxy (nucleus and halo)

frequented by the clusters. They are found

similarly in the halo of M 31 and in nearby

dwarf galaxies. In general they seem to favor

regions which are free of dust. Since it does

not necessarily follow, however, that stars

inhabiting a dust-free region are always the

same kind (i. e., same chemical abundances and

associated patterns of evolution), one cannot

say with certainty whether stars of the globularcluster kind do or do not constitute the predominant population in elliptical galaxies; if

they do, then they are probably the most

numerous kind of stars in the universe. This

is a fundamentally important question on

which more work is needed.

Color-magnitude diagrams

One of the best ways to characterize the stars

that compose a stellar system is to plot their

color indices (or their spectral types) against

165
166 SMITHSONIAN CONTHIBUTIONS TO ASTROPHYSICS

- 6

My

- 4

- 2

0

+2

• 4

+6

+8

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i i i

\

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1 1

MI3

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-.4 +.4 +JB -t-1.2 +1.6

FIGURE 1.—Schematic color-magnitude diagrams for M 3 and M 13. Sequences for the solar neighborhood are sketched with

dashed lines for comparison. The diagrams have been adjusted vertically so that the mean of the cluster-type variables

(circle) in M 3 falls at M,=0 and so that the evolutionary breaks from the main sequence fall at the same bolometric magnitude

(dotted line).
NEW HORIZONS IN ASTRONOMY 167

their magnitudes. In contrast to the diversity

of color-magnitude diagrams found among

open clusters which inhabit the galactic disk,

the color-magnitude diagrams of globular

clusters all have basically similar features.

This does not mean that they are identical.

There are indeed important differences among

them, but the differences mainly pertain to the

positions of the sequences rather than to the

existence or nonexistence of various sequences.

Color-magnitude diagrams for two wellknown globular clusters, M 3 and M 13, are

sketched in figure 1. The color indices span

an extremely broad range; the bluest stars have

indices around —0.5 on the PQ— V system while

the reddest are around +1.5. The brightest

stars in globular clusters have absolute photovisual magnitudes of about —3, while the

faintest are presumably much below the present

limit of observation at the bottoms of the

diagrams shown here.

In the lower part of each diagram, from

M,~+ 4 downward, is a segment of the main

sequence. Stars which formerly lay a little

above M f ~ + 4 on the main sequence have

recently evolved away from it and they

presently form the upper branches of the

diagram. Stars which started from somewhat

brighter levels of the main sequence, say

above M,~+3 , have long since completed

their evolutionary sequences and passed on

to the graveyard. When a star evolves away

from the main sequence it evidently moves

into the subgiant region at the right of the

main sequence and climbs upward into the

red-giant branch. Since the total ascent is

about seven magnitudes, the star is consuming

fuel several hundred times faster during its

red-giant phase than during its childhood on

the main sequence. As a result it doesn't live

very long as a red giant. When it starts to

run short of hydrogen, the star is believed to

move toward the left along the horizontal

branch and then dribble downward at the

extreme left of the diagram to the graveyard.

Its death throes are suspected of being occasional nova outbursts, after which it collapses

into a degenerate white dwarf and cools itself

into oblivion.

The time required for the color-magnitude

diagram to reach its present configuration is

uncertain, but it apparently is of the same

general order as the age of the universe judged

from the recession of galaxies, say, 5 to 10

billion years. Each star probably spends more

than half of its life on or near the main sequence,

where it uses up only 10 to 15 percent of its

hydrogen. As it proceeds up the vertical

branch, its rate of rise is roughly proportional

to the luminosity, and its radius expands as

the square root of the luminosity. By the

time it reaches M,~0, the star is evolving quite

fast. It requires only about 2X108 years to

go the rest of the way to the top of the diagram

and about an equal length of time to traverse

the horizontal branch. The nova stage, if

indeed that is the star's actual route to the

graveyard, must be less than 109 years' duration. Beyond that, its evolutionary schedule

is unknown.

Near the top of the red-giant branch, the

star starts to fluctuate small fractions of a

magnitude at a relatively slow rate, its amplitude increasing as the tip of the branch is

approached. In the same region of the colormagnitude diagram one occasionally finds a

bonafide long-period variable, which suggests

that the fluctuation may build up briefly to

large amplitudes as the hydrogen-spent star

starts to collapse and slide back down the redgiant branch. In any case, it apparently

quiets down again as it travels rapidly toward

the left along the horizontal branch. Then

suddenly it breaks into another fit of oscillation, this time as a cluster-type variable. Its

period and light curve are relatively steady,

but they are believed to undergo gradual

secular changes as the star pursues its evolutionary course through this region of the

diagram. Present data suggest that the period,

starting initially at about one day, becomes

progressively shorter, while the amplitude

rises to a maximum and then falls again. This

fall in amplitude is accompanied by a change

in the shape of the light curve which may be

associated with a change in the mode of pulsation. At the blue edge of the variable-star

region the star quiets down again until it slides

off the blue end of the horizontal branch. If

the star then passes through a nova stage, one
168 SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS VOL.1

might count it as the third fit of instability

along the evolutionary track, but it evidently

differs in kind from the other two.

The foregoing paragraphs outline the general

features of globular-cluster diagrams and their

evolutionary interpretation as we understand

them today. While certain basic facts now

seem to be fairly definite, it should be emphasized that much of this information is only tentative and some of it is even speculative.

Moreover, there are a few items which were

omitted altogether, because we haven't the

foggiest notion how to fit them into the over-all

evolutionary picture. These include the longperiod Cepheids (between M,^ — 1 and

M,~*—5) found in some of the clusters, the

planetary nebula in M 15, and the occasional

blue stars found above the horizontal branch in

a few clusters.

Altogether, there are so many points on

which more photometric material is needed that

it is difficult to suggest what should be undertaken first. In some cases, the question is

largely one of observing more stars or sampling

more clusters. This is true, for example, if one

wishes to understand the nature of red-giant

fluctuation, the properties of the red long-period

variables, or the role of the long-period Cepheids;

it is also true of investigating the differences

between cluster diagrams, which we shall discuss shortly. In other cases, the problem is

rather to improve the accuracy of existing data.

It has sometimes been tempting to overrate the

finality of the present material and to neglect

the need for further investigation.

Among the various branches of the cluster

diagrams, the main sequence is observationally

the most difficult. Recent developments in

photoelectric photometry, however, have made

it possible to measure extremely faint stars with

complete freedom from the scale errors that

have hampered photographic efforts. As a

result of such developments, it has been found

that the stars comprising the main sequence of

a globular cluster do not fall on the ordinary

main sequence with which we are familiar.

They are subdwarfs. This result is of cosmogonic significance because it leaves little

doubt that globular clusters are made from a

different recipe from that of the majority of

stars in the solar neighborhood; a disparity in

age alone would hardly seem sufficient to

account for it.

There are some fascinating differences between the color-magnitude diagrams of various

globular clusters. Features which differ noticeably from cluster to cluster include: (1) the

position of the main sequence, (2) the position

(colorwise) of the top of the red-giant branch,

(3) the shape and tilt of the red-giant branch,

(4) the number of cluster-type variables, (5)

the populousness of the horizontal branch at

the right of the cluster-type variables, and (6)

the shape and tilt of the horizontal branch.

The last four of these features seem to be

mutually related; if one intercompares all of the

color-magnitude diagrams which have been

photoelectrically calibrated on the same color

system, he finds that they can be arranged into

a series (call it series A) in which these last

four features tend to change mutually and

progressively from one diagram to the next.

The number of cluster-type variables thereby

becomes a handy index for categorizing globular

clusters and their diagrams. It is rather puzzling to note, however, that the second feature

in the list above suggests arranging the diagrams into a different series (call it series B),

and that it is this latter arrangement which

spectroscopic observations tend to support.

Data bearing on the first feature in the list

above are still too scanty to tell us which

series, if either, the main sequences will favor,

but a prediction can be ventured: if series A

truly represents the effects of chemical abundances upon evolution, the arrangement of the

main sequences ought to be in accord with it.

Actually, series B has somewhat the appearance of being series A folded in the middle, but

the number of clusters investigated is too

small to provide more than a hint of any such

relationship. If one intercompares a group of

cluster diagrams arranged in order according

to series A, the tops of the red-giant branches

tend to be progressively bluer until one comes

to M 92, after which the red-giant branches

are progressively redder. Both extremes of

series A might thus be associated with the red

extreme of series B, and we could then forget

about series B having any independent significance. This happy unification seems to

run into grief, however, if we attempt to
NEW HORIZONS IN ASTRONOMY 169

interpret abundances from the spectra and

simultaneously to invoke abundances to account

for series A. If abundance differences are

reserved for explaining series A, and if B cannot

be associated with it, then perhaps we should

seek a phenomenon which does not affect the

general evolutionary pattern but which does

affect the swollen atmospheres associated with

the stars during their red-giant phases. The

fact that all the giants of a cluster are affected

to somewhat the same degree suggests the

possibility of an environmental phenomenon,

such as surface accretion during passage of the

cluster through the galactic plane or nucleus,

but such environmental effects have not been

regarded with much favor. A third possibility

for explaining the existence of the two series is

based upon the fact that abundance differences

themselves are not necessarily one-dimensional;

in principle there are as many parameters as

there are elements in the periodic table. We

can imagine, for example, that some of the

features of the color-magnitude diagram (maybe

those abiding by series A) might be sensitive

to the hydrogen/helium ratio, while others

(maybe those abiding by series B) might be

sensitive to the hydrogen/metals ratio. In any

case, the problem of interpreting these series is

a very challenging one.

Photometric observations of high precision

are badly needed in more globular clusters if

the differences between them are to be fully

understood. Although many color-magnitude

diagrams of various sorts exist, only a few have

been calibrated photoelectrically and adjusted

accurately to a standard color system. There

are plenty of favorably situated clusters whose

diagrams would be valuable. Table 1 is a list

of clusters which lie 20° or more from the

galactic plane. Twenty-four of them are in

the Southern Hemisphere while only 13 are in

the Northern Hemisphere; however, several

of the southern clusters can be reached fairly

well by northern telescopes. The modulus

represents a rough estimate of the magnitude

of the horizontal branch; for the upper branches

to be well defined, the color-magnitude diagram

should extend about two magnitudes fainter

than the modulus without serious deterioration

of accuracy. Since integrating techniques now

enable one to obtain photoelectric observations

to about the same limit that can be reached

photographically, most of these clusters can

be reached with telescopes of modest aperture if

they are adequately equipped. The approximate number of known variables, the integrated

spectral class (or alternatively the integrated

color index), and the approximate distance from

the galactic nucleus are all parameters by

which clusters can be grouped and compared.

We have already cited the number of variables

as an index to the form of the color-magnitude

diagram according to series A. We shall return

in a moment to a discussion of the relationship

of these integrated properties to the distance

of the cluster from the galactic nucleus.

Among the clusters in table 1, color-magnitude diagrams suitable for intercomparing

the upper branches have thus far been published for NGC 4147, M 3, M 5, M 13, M 92,

M 2, M 10, and M 15. Work is now in progress

on one or two others, but the majority haven't

been touched. It should also be mentioned

that there are several clusters which are too

interesting to neglect in spite of the fact that

they lie at less favorable galactic latitudes

than those listed in table 1. Outstanding

examples with low moduli are w Centauri, M 4,

M 22, and NGC 6397. Clusters lying close to

the galactic nucleus include NGC 6304, NGC

6316, NGC 6355, NGC 6453, and NGC 6522.

There are also two or three massive galactic

clusters of great age, similar perhaps to M 67,

whose color-magnitude diagrams may help to

clarify evolutionary differences between halo

and disk populations.

It is much more difficult to reach significant

samples of the main sequences than similarly

to reach the upper branches. It requires a

large telescope and a great deal of photoelectric

observing time under first-rate conditions.

Thus far, M 13 is the only cluster for which the

main sequence has been firmly located. Some

tentative photoelectric results were used for

including M 3 in figure 1, and similar material

also now exists for M 15. Earlier appearances

of normal main sequences in the photographically extrapolated diagrams for M 3 and M 92

have turned out to be spurious. It may be

several years before a definitive amount of

data on the relationships of globular-cluster

main sequences can be assembled.
170 SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS

TABLE 1.—Globular clusters for which |6|>20c

Clutter designation a 1950 t 1950

NGC 104 47 Tuc 00 2l!°9 -7 2 21

288 00 50.2 -2 6 52

362 01 00.6 -7 1 07

1261 03 10.9 -5 5 25

1841 04 52. 5 -8 4 05

1851 05 12. 4 -4 0 05

1904 M 79 05 22.2 -2 4 34

2419 07 34.8 +39 00

4147 12 07. 6 +18 49

4590 M 68 12 36.8 -2 6 29

5024 M 53 13 10.5 +18 26

5053 13 13. & +17 57

5272 M 3 13 39.9 +28 38

5466 14 03. 2 +28 46

5634 14 27.0 -0 5 45

5694 14 36.7 -2 6 19

5824 15 00.9 -3 2 53

5897 15 14. 5 -2 0 50

5904 M 5 15 16.0 +02 16

6171 M107 16 29.7 -1 2 57

6205 M 13 16 39. 9 +36 33

6218 M 12 16 44 6 -0 1 52

6229 16 45.6 +47 37

6254 M 10 16 54.5 -0 4 02

6341 M 92 17 15.6 +43 12

6684 18 44. 1 -6 5 14

6752 19 06.4 -6 0 04

6809 M 55 19 36.9 -3 1 03

6864 M 75 20 03.2 -2 2 04

6934 20 31. 7 +07 14

6981 M 72 20 50.7 -1 2 44

7006 20 59. 1 +16 00

7078 M 15 21 27.6 +11 57

7089 M 2 21 30.9 -0 1 03

7099 M 30 2137.5 -2 3 25

NGC7492 23 05.7 -1 5 54

IC4499 14 52.7 -8 2 02

Estimated modulus

(M)

mag

14. 5 15. 7

15. 7

16. 3

19.4

17.0

16. 1

16.7

16. 4 15.7

16. 4 17. 1

17.7

16.2

15.2

16.2

14. 8

15.2

17. 7

15.2

15.2

14.2

14. 5

18. 1

16.9

17.0

18.8

15. 9

16. 2

15.5

17.9

Number of variables

8

2

14

3

5

36

4

28

42

10

186

18

7

0

97

24

15

1

21

2

16

1

2

11

51

31

20

66

17

3

9

Integrated spectrum

F5

F3

F5

A5

A6

A8

F2

F4

A9

F5

F7

G2

F2

F7

F6

GO

A5

Gl

F9

G2

Fl

FO

FO

A7

Distance from

gal. nucleus

kpc

8

15

12

20

65

24

13

19

17

13

17

19

22

8

7

6

9

4

28

4

10

5

4

25

15

16

39

12

13

9

32

The foregoing discussion of color-magnitude

diagrams is presented in terms of ordinary

two-color photometry. Much can be added,

however, by extending observations to three or

more colors. In particular, the addition of

ultraviolet enables one to compare the relative

strengths of the Balmer regions in various

stars. This can be done for instance by

plotting a diagram of U-B against B-V.

Results in M 13 show that the red giants and

subgiants have an ultraviolet excess of about

0.05 magnitude, while the very blue stars on

the horizontal branch have very pronounced

ultraviolet deficiencies. Rather oddly, there

is a lonesome B 2 star lying about two magnitudes above the horizontal branch in M 13

which does not exhibit this ultraviolet deficiency. In M3 and NGC 4147, the ultraviolet

excess of the red giants and subgiants is several

times greater than in M 13. Similar information for other clusters would be very valuable;

it might provide an important clue to the reason

for the differences between their color-magnitude diagrams.

Multicolor photometry and spectrophotometry have not yet been applied to individual

stars in globular clusters, and the results of

any multicolor exploration are likely to pay

handsome rewards. Although the useful range

of a photometer becomes limited to brighter

objects when the color bands are made narrower,

the bright stars in globular clusters are still

within easy reach of conventional six-color

observations with a telescope of modest aperture. On the other hand, narrow-band spectrophotometry will be quite difficult.

Spectra

Spectroscopically, the individual stars in globular clusters are rather difficult objects; hardly

any of them are brighter than 14th magnitude

in the blue region. Until recently, the spectra

have consequently been of quite low dispersion.

These have been suitable, however, for spectral
NEW HORIZONS IN ASTRONOMY 171

classification and for detection of the grosser

peculiarities of the cluster stars. It was already

recognized a decade ago, for example, that

cyanogen absorption is abnormally weak and

that the hydrogen lines are a little more enhanced than in normal stars of the same

absolute magnitude.

During the last few 3rears, the fast new

coud6 cameras at Palomar have made it possible to reach some of the brighter cluster stars

with a dispersion of 38 A/mm and, in two

cases, of 18 A/mm. Because of the very large

collimator-to-camera ratios for these dispersions, the slit of the Palomar spectrograph can

be made wide enough to receive nearly all of a

star image during good "seeing" without causing

a loss of photographic resolution in the spectrum. The first four such spectra at 38 A/mm

were obtained in 1952 with the idea of checking spectroscopically on the photometrically

observed difference between the red-giant

branches of M3 and M92, which at that time

were the only two clusters whose color-magnitude diagrams could be intercompared on a

standard system. These spectra not only confirmed the reality of an intrinsic difference

between these two clusters but also showed

quite dramatically that the red giants in M92

are very extraordinary objects. They are

markedly different from any other known stars.

Additional spectra of the same dispersion were

soon obtained in M92 and, subsequently, in

most of the globular clusters for which intercomparable color-magnitude diagrams now

exist as well as in several others not yet photometrically investigated.

The most striking peculiarity of the redgiant spectra in some of the clusters is the

extreme weakness of the metallic lines. Superficially, these spectra appear to be of very

much earlier types than the color indices of the

stars suggest. In M 92, for instance, the

metallic lines seem to have roughly the strengths

usually found in F-type stars, whereas the color

indices would normally be associated with

early K stars. This disparity largely vanishes,

however, if the spectral classification is based

upon the ratios of certain metallic lines instead

of^upon their absolute strengths. In other

words, the excitation temperature is much less

discordant with the color temperature if it is

inferred, as it should be, from the ratios of line

strengths instead of from the absolute strengths

themselves.

A few particular ratios deserve mention. As

the excitation temperature increases, the ratio

of X4254 to X4250 should decrease, whereas the

ratio of X4247 to X4250 should increase. The

feature at 4254 A is a resonance line of neutral

chromium (0.0 volt excitation potential), the

feature at 4247 A is a blend of scandium II with

a high excitation line of neutral iron (3.3

volts), and the comparison line at4250A arises

principally from a low level of neutral iron

(1.6 volts). Another useful ratio which decreases with increasing temperature is Cr I

4275 to Fe I 4271, and another ratio which

increases with increasing temperature is Ti II

4394 to Fe I 4405.

M 92 is apparently an extreme case of the

peculiarities cited above, although M 15 runs

it a close second. In some of the other clusters,

such as M 3 and M 13, these effects are relatively

mild. Within any one cluster, the spectra all

show roughly the same degree of peculiarity.

In general, this degree of peculiarity appears to

be correlated with the redness of the red-giant

branch (i. e., with series B) in the sense that

those clusters whose red-giant branches lie

the farthest toward the right in the colormagnitude plane are the least peculiar spectroscopically.

All of the spectroscopic evidence suggests

that the weakness of the metallic lines in some

of the cluster spectra is a bona fide abundance

phenomenon. Quantitative abundance estimates have been made on the basis of spectra

obtained at 18 A/nun in M 92 and M 13, which

represent opposite extremes. Equivalent

widths of suitable metal lines in an M 92

spectrum are found to be about 60 percent of

those in an M 13 spectrum, implying that the

metals in M 92 are at least a factor 10 less

abundant than in M 13. Line ratios in these

same spectra assure us that the excitation

temperatures in M 92 and M 13 are about

the same in spite of the difference in the color

indices; perhaps the color difference is due to

the fact that metallic lines chop different

amounts out of otherwise similar continua.

As pointed out earlier, however, the use of

abundances to account for spectroscopic differ-
172 SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS

ences (and therefore series B) poses a dilemma,

because we also feel obliged to call upon

abundances to explain the various evolutionary

features which behave according to series A.

Possible ways around this dilemma, such as

allowing more than just the hydrogen/metals

ratio to vary, have already been mentioned.

Integrated properties

In addition to the classification of colormagnitude diagrams, there are several other

informative parameters which pertain to each

cluster as a system or as a unit. These include

integrated luminosity or magnitude, integrated

color index, integrated spectral-energy distribution, integrated spectral class, mass, massluminosity ratio, distance, reddening, position

in the galaxy, motion, spatial profile, and so on.

Since the integrated luminosity and color

represent the summed effects of all of the

members of a cluster, one should in principle

be able to subtract the contributions of the

resolved stars from the total observed luminosity, and thereby deduce how much light is

contributed by members which are too faint

to be individually resolved by present telescopes,

say, fainter than mPf=23. In practice, however, uncertainties in the luminosity function

(i. e., in the numbers of stars per magnitude

interval) are much too large and the quantity

sought is much too small for such a procedure

to succeed. Direct measurements of this

unresolved background by itself appear to show

some indications of a substantial white-dwarf

population, but the uncertainties are presently

too large to justify much confidence.

As one might expect from the differences

between color-magnitude diagrams, the fractions of the total light contributed by the

various branches differ somewhat from cluster

to cluster. Typically, about half of the visual

light comes from the red giants above Af,=0,

perhaps 20 percent each from the horizontal

branch and from the subgiants, and the

remainder from the main sequence downward.

In view of the broad range of color indices in

each cluster diagram, it is not surprising that

the integrated color indices of globular clusters

are not identical. One cannot correctly determine the reddening of an individual cluster

simply by comparing its color index with an

unreddened mean.

The integrated spectral-energy distribution of

a stellar system such as a globular cluster is, of

course, quite different from that of a single star

having the same color index. This means that,

regardless of data-reduction procedures, the

ordinary color index of a cluster is rather sensitive to the actual color response of the photometer used.

The spectral energy distribution of a globular

cluster is also very different from that of an

elliptical galaxy, indicating that the colormagnitude diagrams must be fundamentally

unlike unless the distribution of stars in the unevolved region of the cluster diagram is radically

different from that for an elliptical. If the diagrams are assumed to be of the same form, one

can show that the ratio of the number of K stars

to the number of F stars on the main sequence

of an elliptical galaxy would need to be literally

hundreds of times greater than on the main sequence of a globular cluster in order to account

for the observed disparity in their integrated

spectral-energy distributions. Since the present

data can be fitted equally well by assuming unlike diagrams, the identification of the predominant population in ellipticals remains an

open question. The importance of further work

on this problem has already been stressed.

For much the same reason that the color index

is sensitive to the wavelength bands used for the

observations, the estimation of the integrated

spectral class of a stellar system such as a globular cluster is dependent upon the spectral region utilized. The blue stars on the horizontal

branch play a bigger role at short wavelengths

than at long wavelengths. Integrated spectra

of 50 clusters have been obtained by throwing

the clusters sufficiently far out of focus to smear

a blend of many stars uniformly over the spectrograph slit. The spectral types listed in table

1 were derived from integrated spectra obtained

in this manner. Some of these spectra are currently being reexamined in an effort to take

better account of the peculiarities now recognized; the present tendency is to assign them

later spectral types than were formerly estimated.

An attempt has been made to determine the
NEW HORIZONS EN ASTRONOMY 173

mass of M 92 from the dispersion in radial velocities among red giants in its nucleus. This

was done by applying the virial theorem to

velocities measured on 23 Palomar spectrograms of 15 stars. Some of the red giants for

which two spectrograms were obtained show

evidence of intrinsic velocity variations, which

is not surprising in view of the photometric

variations already described. As a result of

these variations, the uncertainty in the true

velocity dispersion is disconcertingly large.

Nevertheless, with some recent refinements in

theory the results yield an estimated mass of

1.4 X 10s solar units, the upper limit being perhaps twice that value. The corresponding

mass-luminosity ratio is of the order of unity,

which is several times smaller than the average

for our own galaxy as a whole and is about a

hundred times smaller than that of an elliptical

galaxy. The interpretation of these comparative ratios in terms of differing populations,

luminosity functions, and nonluminous content4'

are not yet clear.

The distances of globular clusters are conventionally estimated by assigning the clustertype variables a mean absolute magnitude of

zero, or by comparing the brightest stars with

those belonging to other clusters whose variables

have been calibrated. It should be remembered

that the assignment of zero absolute magnitude

to the cluster-type variables is only a rough approximation which should not be taken too

seriously when color-magnitude diagrams are

intercompared. In figure 1, for example, the

cluster-type variables are a half magnitude apart

because the diagrams were adjusted to place the

evolutionary breaks from the main sequences at

the same bolometric magnitude.

Since most globular clusters inhabit regions

well outside the absorbing material in the

galactic disk, the photometric moduli (blue)

should, on the average, be reduced by 0.25

esc b in order to estimate true distances. While

a statistical correction of this sort is useful

when one is dealing with large-scale problems

involving many clusters, it does not provide a

reliable estimate of absorption for individual

clusters due to the probable patchiness of the

absorbing blanket in which we are imbedded.

Unfortunately, other criteria of reddening and

absorption such as comparison of U, B, V data

with the solar neighborhood or matching of

colors of the cluster-type variables are also open

to question. Consequently, the existence of

accurate photometric data would not by itself

guarantee a precise distance for a cluster even

if the absolute magnitude of the cluster-type

variables were exactly known.

The fact that globular clusters differ significantly from one another in their diagrams and

in their over-all characteristics suggests that

they probably were not all formed under identical circumstances. It is therefore of interest

to consider what clues we have as to their times

and places of origin during early epochs of the

galaxy. When the integrated characteristics

of globular clusters are examined in terms of

their positions in the galaxy, some fascinating

correlations stand out immediately. In particular, the number of cluster-type variables,

which is an index for classification according to

series A, appears to have a marked correlation

with the distance from the galactic plane. If

we examine all clusters which have been

searched for variables, which are not badly

obscured, and whose distances are fairly well

known, we find that the variable-poor clusters

(say, those with 0 to 3 variables) favor regions

much closer to the galactic plane than the

others, hinting that they may have been formed

during a slightly later stage in the dynamic

relaxation of the galaxy. The sample of clusters in table 1 is not well suited for illustrating

this effect because of its cutoff at 20° latitude,

but the effect can still be seen; the mean distance from the galactic nucleus to the variablepoor clusters listed in the table is 9 kpc, while

the mean distance for the variable-rich clusters

is 19 kpc.

It follows, of course, that the various characteristics which are correlated with the number

of variables tend to show this same dependence

on galactic distribution. For example, the integrated spectral type, which is influenced by

the relative shape and populousness of the

upper branches of the cluster diagram, reveals

an effect of this kind. Among the clusters

listed in table 1, those possessing later spectral

types (say, F 7 to G 2) have a mean distance

of 11 kpc from the nucleus, while those pos-
174 SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS

sessing earlier types (say, A 5 to F 6) have a

mean distance of 21 kpc from the nucleus.

Again, the effect would be more distinct if

clusters in the 10° to 20° zone were added to

the sample.

If the centroid of the system of globular

clusters in our galaxy is assumed to coincide

with the galactic nucleus, one can check the

zero point of the cluster-type variables against

the present estimate of the sun's distance from

the nucleus (about 8 kpc). The variable-rich

clusters seem to be quite well centered on the

nucleus if the conventional value of M=0

is adopted for the variables, but the variablepoor clusters appear to prefer a slightly brighter

value, perhaps half a magnitude. While

scarcely more than qualitative due to the

smallness of the sample, this result is nicely

consistent with figure 1.

There are several observable globular clusters

inhabiting the immediate vicinity of the

galactic nucleus. One or two of these may be

transients just passing through, but others are

suspected of being permanent residents and of

being quite peculiar. It would be especially

interesting to see how they fit into the over-all

type-distribution picture described above.

The individual space motions of globular

clusters in the galaxy are unknown because

they are too distant to have reliable proper

motions. Nevertheless, some general conclusions can be drawn from their radial velocities,

which range from —360 km/sec to +290

km/sec. As already remarked, the globular

clusters probably do not share to any large

degree in the general rotation of the galaxy.

The present evidence suggests that the rotational velocity of the cluster system is less than

a third of the rotational velocity of the disk.

It may be difficult to pin this figure down more

closely due to the limited number of clusters

and to their large individual motions. Relative to the clusters, the sun's orbital velocity

is about 200 km/sec in roughly the same direction derived from differential rotation. The

cluster system shows no indication of a general

contraction or expansion.

Efforts have been made to determine whether

globular clusters favor orbits of high ellipticity

passing through the galactic nucleus or whether

they travel in approximately circular orbits at

constant radii from the nucleus. Although

arguments have been advanced for both models,

the present radial velocity data do not make a

clear case for either. Probably the truth is

somewhere in between, and perhaps the situation is further complicated by the form of the

potential function of the galaxy. In any

case, the time required for a cluster to make

one loop or one oscillation is of the same

general order as the period of galactic rotation;

hence, each cluster has passed through the

galactic plane quite a few times since its birth.

The spatial distribution of the stars in a

globular cluster has been the subject of several

observational and theoretical studies. Observationally, the data consist of star counts or surface brightness measures as a function of radius

from the cluster center. To a first approximation, these quantities vary as (constant+r2

)""2,

and the corresponding integrated luminosity

rapidly approaches a well-defined asymptotic

limit for large r. This two-dimensional profile

is very different from that of an elliptical

galaxy, whose integrated luminosity continues

to increase as log (constant+r) and which

shows no tendency for a limit except that

imposed by the domain of its neighbors.

Since the profile of a globular cluster is a

manifestation of internal dynamic conditions,

it is related to the numbers of stars having

various masses, to their mass-luminosity ratios,

and to the manner in which energy is partitioned. We know that nearly all of the light

of a globular cluster comes from stars above its

main sequence and that their masses should all

be of the same order because they all evolved

from the same region of the main sequence;

thus a single mass-luminosity ratio applies fairly

well to the surface-brightness profile. This has

been checked by noting that the difference in

integrated color between inner and outer regions

of a cluster is generally small and that the radial

spread of the red-giant population does not

differ significantly from the spread of the subgiants and the blue stars. However, since the

exchange of energy in the central region of a

cluster is believed to be sufficient for equipartition, one should expect the profile for the

main-sequence membership to be broader than
NEW HORIZONS IN ASTRONOMY 175

the profile for the brighter stars, and there is

observational evidence in support of this.

Differing amounts of internal energy per unit

mass from cluster to cluster probably account

for their differing degrees of over-all compactness.

The origin of globular clusters is still shrouded

in mystery. It seems fairly evident that most of

them in our own galaxy were formed during an

early epoch in its evolution. We do not know,

however, whether they condensed out of interstellar material originally in the domain they

now inhabit or whether they were created in

denser central regions and ejected due to the

high turbulent energy which may have existed

at that time. M 87, which is a very massive

elliptical in the Virgo cluster of galaxies, provides a possible clue to this question. It has

two unusual features: there are an astonishing

number of globular clusters associated with it,

and there is a peculiar blue jet streaming from

its nucleus with several pronounced globules

along it. Is it possible that M 87, because of

its unusual mass and energy, has a slow rate of

dynamic development and is still in the stage

of making globular clusters? Perhaps the

range of integrated color indices of the M 87

clusters will help to decide whether any of them

can be young, but the problem is observationally difficult. It is also of interest to note

that the radial distribution of the globular

clusters in M 87 is approximately the same as

the distribution of the nonclustered population

(judged from the surface brightness profile).

Since the relaxation time of a galaxy is theoretically long compared with the average age of

these clusters, this result implies that the conditions of formation have been such as to produce

similar distributions since the beginning.

In closing, we should not neglect to consider

the value of globular clusters as tools in other

astronomical problems. Their usefulness in

surveying the structure of our own galaxy is

already evident from earlier remarks, but their

importance in extragalactic studies also deserves

comment. Globular clusters belonging to other

galaxies are convenient distance indicators.

Their luminosities can be used for intercomparing the distances of galaxies as far as the

Virgo cluster of galaxies, which serves as a

cosmic yardstick for everything beyond. A

recent estimate based on the globular clusters

associated with M 87 places the Virgo cluster

of galaxies at 10 megaparsecs. If new results

for more distant clusters of galaxies are adjusted

to this scale value, one obtains a Hubble constant of 150 km/sec per megaparsec, whose

inverse is 6.5 billion years. Much more work

is needed on this problem, not only in connection with the photometry of globular clusters

in remote galaxies but also in connection with

establishing a better luminosity function by

which the clusters belonging to different galaxies

can be compared. We must also worry about

whether the luminosity function for clusters is

necessarily the same in all kinds of galaxies.

A fair amount of observational work has been

done on the globular clusters in the Andromeda

galaxy (M 31). In some respects they provide

a better sample than those in our own galaxy

for investigating colors, luminosities, compactnesses, motions, and distribution. The brightest

are around 15th apparent magnitude, which

corresponds to M ~ — 9. They are close enough

that we can even compare their profiles with

those of clusters in our own galaxy to obtain an

obscuration-free check on the distance of M 31.

The present paper outlines some of the

problems in a very broad field. So many different workers have had a part in building our

present state of knowledge that it would have

been foolish to attempt the incorporation of

anything approaching a complete system of

bibliographical references into the present text.

An excellent 80-page bibliography, by Helen B.

Sawyer, already exists (Publ. D. Dunlap Obs.,

vol. 1, p. 383, 1947), but anyone actually planning to work in the field will also want to

become acquainted with at least part of the

work done since that date. Recent publications and current programs include work by

the following: Arp, Artiukhina, Baade, Baum,

W. Becker, Belserene, Bernstein, Cuffey,

Deutsch, Eggen, Fatchikhin, Florya, Gaposchkin, Gascoigne, Greenbaum, Greenstein, Hoag,

von Hoerner, Huang, Idlis, Irwin, H. L. Johnson, Kholopov, I. King, Kron, Kukarkin,

Lohmann, Mayall, Melbourn, Minin, Morgan,

Munch, Oort, Oosterhoff, Osterbrock, Parenago,
176 SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS TOUI

Reddish, Roberts, Roman, Rosino, Sandage, of these activities are summarized and referSavedoff, Sawyer-Hogg, Schurer, Schwarzschild, enced in the 1948, 1952, and 1955 reports of

Shapley, Stebbins, Swope, Tayler, Walker, Commission 37 of the International AstronomWahlquist, Wallerstein, Whitford, Wilkins, ical Union (see Trans. Int. Astron. Union,

O. C. Wilson, van Woerden, and others. Some vols. 7-9).
Stellar Structure and Evolution

By M. Schwarzschild

 1

The theory of the internal constitution of the

stars entered a basically new phase in 1938

with the discovery of the particular nuclear

processes that provide the main energy sources

for the stars, and, soon thereafter, with the

fairly accurate determination of the reaction

rates of these processes.

It had long been known that a star of given

mass and chemical composition has a unique

equilibrium structure. But it was not possible

actually to determine this unique structure

as long as one of the fundamental relations,

that for the rate of energy production by

nuclear processes, was unknown. Now that

the necessary nuclear data are available we

can not only determine the internal structure

of a star in its initial state when it just starts

burning its hydrogen fuel, but we can also

follow, step by step, in a quantitative nonspeculative manner, the evolutionary changes

caused by the nuclear transmutations.

Many possibilities are now within our grasp:

to interpret the Hertzsprung-Russell diagram

in terms of stellar evolution, to determine the

ages of star clusters and even of individual

stars, to sort out the oldest stars and by analyzing them to gain insight into the composition

of our galaxy at its earliest phases, and to

obtain clues to the origin of the elements by

a comparative analysis of the oldest and the

youngest stars.

Only the very first steps have as yet been

taken in this large field of new and fundamental

problems. The future speed and success of

the developments will depend not only on the

number of specialists working in the subject

and on their ingenuity but also, much more

than in earlier decades, on the advances in

neighboring scientific and technical areas.

Clearly, more highly accurate determinations

of reaction rates for a larger variety of nuclear

1

 Princeton University Observatory, Princeton, N. J.

processes will be needed from nuclear physics.

Similarly, for many star clusters color-magnitude diagrams of high precision at faint magnitudes are needed from observational astronomy. In both these fields, which are fundamental to the theory of stellar evolution,

progress has been very rapid and encouraging

in recent years. There are, however, three

other neighboring fields which should be

actively cultivated if the investigation of stellar

evolution in its broadest sense is to progress.

Computing facilities are our first concern.

The basic problem of stellar structure and evolution can be formulated mathematically in

terms of a fourth-order boundary-value problem, whose solution depends on the stellar

mass and chemical composition and varies with

time. To find the solution, numerical methods

have to be employed. Much useful preliminary computation can be done with desk machines. The numerical work needed, however,

to obtain a sufficient range of solutions with

satisfactory detail and accuracy is so extensive

that it cannot be tackled effectively with desk

calculating machines; large electronic machines

are needed. At a small number of institutes

such large machines have generously been made

available for pure astronomical research. Astrophysical research absolutely requires the help

of large electronic computers if it is to grow

from its present promising beginning to the

maturity the problems demand.

The second area we need to develop is photometric spectroscopy with high dispersion to determine the chemical composition of many sample stars. This observational work, which is

basic to many branches of astrophysics but

which is very tedious and time consuming with

traditional techniques, has progressed very

slowly in recent years despite its importance.

These investigations may expand as soon as it

becomes practicable to use electronic image

177
178 SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS

tubes for this purpose. The advent of this new

instrumentation, however, will have its full revolutionizing effect only if at the same time a

larger number of technicians becomes available

to build and maintain the new instruments and

to reduce the data.

The third field is that of the turbulent phenomena in stellar atmospheres. Recent investigations have shown that these phenomena are

of interest not only for the theory of stellar atmospheres but also play a decisive role in determinations of the stellar radius for many types

of stars; hence they are essential for the complete theory of stellar structure. Laboratory experiments and theoretical investigations in hydrodynamics have resulted in great progress,

during the last decade, in our understanding of

turbulence. However, in this difficult subject

extrapolation from the laboratory to the stars

will always involve some risk; therefore, direct

observations of turbulence in stellar atmospheres would be highly valuable. This is, in

principle, possible for the photosphere of the

sun. In practice, however, it is extremely difficult because of interference from the disturbances of the earth's atmosphere. In consequence, very little progress has been made in

this field during the past 50 years. Now, however, observations from high-altitude vehicles

appear to open up new possibilities; if they can

be realized, the final major difficulty in the

theory of the stellar interior may be resolved.

At present, no serious block seems probable in

the necessary development of any of these fields.

Thus, research in the theory of stellar structure

and evolution presents many unsolved, but solvable, basic problems and may result in a great

advance in answering objectively certain cosmological questions. The rate of successful progress in this field will depend essentially on the

ingenuity and the number of workers who apply

themselves to it.
Radio Sources

By John D. Kraus *

Perhaps you have seen a distant landscape

compared in two photographs, one taken in

ordinary light and the other in infrared. In the

ordinary photograph the distant objects are

largely hidden by haze, but in the infrared

picture the distant hills and valleys stand out

clearly. Infrared rays are longer than light

rays and more easily penetrate the haze.

Radio waves are millions of times longer than

light and have an even greater power of

penetration than infrared. Because of this

fact, radio—applied to astronomy—cuts

through the cosmic haze caused by the clouds

of gas and dust in interstellar space and

presents a view of vast regions that will always

remain hidden to ordinary telescopes. Radio

has, in effect, opened a new window on the

universe.

It was about 1873 that Maxwell postulated

the possibility of radio waves. About 15

years later Hertz produced them and demonstrated their transmission over a distance of a

few feet. In 1901 Marconi received radio

signals across the Atlantic and an era of worldwide radio communication was begun. Thirty

years later Jansky detected radio waves from

the center of our galaxy and initiated a new

method for exploring the cosmic universe. It

was 60 years from Maxwell to Jansky, from a

man's idea to its cosmic application. And

today, some 20 years after Jansky's discovery,

this application of radio to astronomy is

rapidly developing into a new science called

radio astronomy.

Radio observations of outer space are made

with an antenna and a radio receiver. The

antenna gathers and focuses the radio waves as

does the lens of an optical telescope, while the

receiver detects and records the signal much as

does the photographic plate of an optical

1

 Ohio State University, Columbus, Ohio.

telescope. By analogy, the antenna and receiver may be called a radio telescope.

As observed with a radio telescope, the sky

looks strangely different. None of the familiar

stars of the night sky are observable, but in

their places are many radio sources forming

totally unfamiliar constellations. Some of these

sources can be identified with supernova remnants and other gaseous nebulosities and with

extragalactic nebulae. There are a couple

dozen sources thus identified, at least tentatively, but there are hundreds more not

identified with any optical object and the bulk

of them may never be. It is as though radio

and optical telescopes "see" two different

universes with only an occasional object

observable by both. And those objects that

can be detected by both appear to be of different

sizes and shapes in the two telescopes. For

example, the sun has a visual diameter of about

one-half degree, but at certain radio wavelengths the diameter measures several degrees

because the radio waves originate in the upper

regions of the corona or solar atmosphere.

If our eyes were sensitive to radio waves

instead of to light, the winter sky would appear

somewhat as suggested in figure 1. The

contours are lines of equal surface brightness

showing the radio background. The background is a maximum along a ridge close to the

plane of our galaxy or Milky Way system. On

this background are dots and circles that show

the position of some of the strongest or brightest

radio sources. Their brightness is indicated by

a radio magnitude scale analogous to the visual

magnitude scale. Some of the sources are of an

extended type with diameters of more than 1°.

These are shown by the open circles. Others

are more localized or pointlike, and these are

indicated by the solid dots. A few sources

have been identified with optical objects and

179
SMITHSONIAN CONTRIBUTIONS TO ASTROPHYSICS

•TV

ROSETTE NEBULA

N6C2244// /

- 4 0

FIGURE 1.—Map of the radio sky at a wavelength of 124 cm made with the radio telescope at Ohio State University

(KoandKraus, 19SS).

for these the optical designation is given. Note

for example the Crab Nebula (Ml), which is a

supernova remnant, or the gaseous nebulae

designated IC 443 and NGC 2244. However,

many of the sources on the map have not been

identified with an optically observable object

and this is true for most radio sources, particularly the weaker ones not shown on the map.

The identification of radio sources with

optical objects is a challenging problem that

requires close cooperation between radio and

optical astronomers. Accurate radio position

and size determinations are vital in this work.

One of the outstanding achievements of radio

astronomy was the identification of the second

brightest radio source (other than the sun)

with a faint pair of colliding galaxies at a distance of some 200 million light-years. Teamwork by radio astronomers at Cambridge,

England, and Sydney, Australia, and optical

astronomers at Mount Palomar was required

for several years before the identification was

made (Baade and Minkowski, 1954). If this

object were at 10 times the distance it would

still be a readily observable radio source but

would then be at or beyond the range of the
NEW HORIZONS IN ASTRONOMY 181

largest optical telescope. It is this fact which

suggests that radio telescopes may open to exploration vast regions of space that lie beyond

the range of any optical instruments.

The known radio sources are now numbered

in the hundreds, with lists published by Ryle,

Mills, Bolton, Brown, Haddock, and Ohio

State University. However, those that are

reliably established are much less numerous

and those identified with optical objects are

only a dozen or two (Pawsey, 1955).

The work of radio exploration has barely

begun and great discoveries lie ahead as larger

radio telescopes are built and more radio sources

are found. Many of these sources may emit

too little light for optical detection or may be

hidden by clouds of cosmic dust, while many

others may be at distances beyond the range

of any optical telescope. Studies of these

sources, individually and collectively, are bound

to supplement and alter our purely optical

picture of the universe in a most significant

way. From Maxwell to Hertz, to Marconi, to

Jansky, to the outermost reaches of our universe

are big steps, but they do not end there and

new advances are being made daily. Perhaps

you may become the radio astronomer who

will make the next big step.

References

BAADE, W., AND MINKOWSKI, R.

1954. Astrophys. Journ., vol. 119, p. 206.

Ko, H. C , AND KBAUS, J. D.

1955. Ohio State Univ. Radio Obs. Rep. No. 4,

June 22; in Sky and Telescope, vol. 14,

p. 370.

PAWSEY, J. L.

1955. Astrophys. Journ., vol. 121, p. 1.

o