Milky Way, and one portion of the Galaxy—between Aquila, Lyra, and
Cygnus—was called by Birmingham “the red region in Cygnus.” Yellow and
orange stars seem to be most abundant in the constellations, Cetus,
Pisces, Hydra, and Virgo, and the white stars in Orion, Cassiopeia, and
Lyra.
CHAPTER III.
THE DISTANCES AND MOTIONS OF THE STARS.
The determination of the distances of the stars from the earth has
always formed a subject of great interest to astronomers. The earlier
observers appear to have thought that the problem was an insoluble one.
The famous Kepler, judging from what he called the “harmony of
relations,” came to the conclusion that the distance of the fixed stars
should be about 2,000 times the distance of Saturn from the sun. Saturn
was then the outermost planet of the solar system. The distance of even
the nearest star, as now known, is about 14 times greater than that
supposed by Kepler. Huygens thought the determination of stellar
distance by observation to be impossible, but made an attempt at a
solution of the problem by a photometric comparison between Sirius and
the sun. By this method, he found that Sirius is probably about 28,000
times the sun’s distance from the earth, but modern measures show that
this estimate is far too small, the distance of Sirius being probably
over 500,000 times the sun’s distance, or about 18 times greater than
Huygens made it.
When the Copernican theory of the earth’s motion round the sun was first
advanced, it was objected that, if the earth moved in a large orbit, its
real change of place should produce an _apparent_ change of position in
the stars nearest to the earth, causing them to shift their relative
position with reference to more distant stars. Copernicus replied to
this objection—and we now know that his reply was correct—by saying that
the distance of even the nearest stars was so great that the earth’s
motion would have no perceptible effect in changing their apparent
position in the heavens; in other words, the diameter of the earth’s
orbit round the sun would be almost a vanishing point if viewed from the
distance of the nearest stars. This explanation of Copernicus was at
first ridiculed, and even the famous astronomer, Tycho Brahé, could not
accept such a startling conclusion. This celebrated observer failed
indeed to detect by his own observations any annual change of place in
the stars, but he fancied that the brightest stars showed a perceptible
disc, like the planets, a fact which, if true, would imply that, if the
distance of the stars was so great as Copernicus supposed, their real
diameter must be enormous. The invention of the telescope, however,
dispelled this delusion of Tycho Brahé, and showed that even the
brightest stars showed no perceptible disc. This was proved by Horrocks
and Crabtree, who noticed that, in occultations of stars by the moon,
the stars disappeared instantaneously, a fact which proved that the
apparent diameter of the stars must be a very small fraction of a second
of arc.
Galileo suggested that possibly the distance of the nearer stars might
be determined by careful measures of double stars, on the assumption
that the brighter star of the pair—if the difference in brilliancy is
considerable—is nearer the earth than the fainter star. He says (in his
“_Opere di Galileo Galilei_”), “I do not believe that all the stars are
scattered over a spherical superficies _at equal distances from a common
centre_, but I am of opinion that their distances from us are so various
that some of them may be two or three times as remote as others, so that
when some minute star is discovered by the telescope close to one of the
larger, and yet the former is highest, it may be that some sensible
change might take place among them.” Acting on this idea, Sir William
Herschel, at the close of the eighteenth century, made a careful series
of measures of certain double stars. He did not, however, succeed in his
attempt, as his instruments were not sufficiently accurate for such an
investigation, but his labours were rewarded by the great discovery of
binary or revolving double stars, most interesting objects, which will
be considered in the next chapter.
Numerous but unsuccessful attempts were made by Hooke, Flamsteed,
Cassini, Molyneux, and Bradley, to find the distance of some of the
stars. Hooke, in the year 1669, thought he had detected a parallax of 27
to 30 seconds arc in the star Gamma Draconis, but we now know that no
star in the heavens has anything like so large a parallax. It must be
here explained that to find the distance of any star from the earth, we
must first measure its “parallax,” which is the apparent change in its
place due to the earth’s motion round the sun. As the earth makes half a
revolution in six months, and as the earth’s mean distance from the
sun—or the radius of the earth’s orbit—is about 93 millions of miles,
the earth is, at any given time, about 186 millions of miles distant
from the point in its orbit which it occupied six months previously. The
apparent change of position in a star’s place, known as parallax, is
_one-half_ the total displacement of the star as seen from opposite
points of the earth’s orbit. In other words, it is the angle subtended
at the star by the sun’s mean distance from the earth. The measured
parallax of a star may be either “absolute” or “relative.” An “absolute
parallax” is the actual parallax. A “relative parallax” is the parallax
with reference to a faint star situated near a brighter star, the faint
star being assumed to lie, as suggested by Galileo, at a much greater
distance from the earth. As, however, the faint star may have a small
parallax of its own, the “relative parallax” is the difference between
the parallaxes of the two stars. Indeed, in some cases a “negative
parallax” has been found, which, if not due to errors of observation,
would imply that the faint star is actually the nearer of the two. From
the observed parallax, the star’s distance in miles may be found by
simply multiplying 93 millions of miles by 206,265 and dividing the
result by the parallax. To find the time that light would take to reach
us from the star—the light journey as it is called—it is only necessary
to divide the number 3·258 by the parallax.
In attempting to verify the result found by Hooke for the parallax of
Gamma Draconis, Molyneux and Bradley found an apparent parallax of about
20 seconds of arc, thus apparently confirming Hooke’s result, but
observations of other stars showing a similar result, Bradley came to
the conclusion that the apparent change of position was not really due
to parallax, but was caused by a phenomenon now known as the “aberration
of light,” an apparent displacement in the positions of the stars, due
to the effect of the earth’s motion in its orbit round the sun combined
with the progressive motion of light. The result is that “a star is
displaced by aberration along a great circle, joining its true place to
the point on the celestial sphere towards which the earth is moving.”
The amount of aberration is a maximum for stars lying in a direction at
right angles to that of the earth’s motion. The existence of aberration
is an absolute proof that the earth does revolve round the sun, for were
the earth at rest—as some paradoxes contend—there would be no aberration
of the stars. This effect of aberration must, of course, be carefully
allowed for in all measures of stellar parallax. To show that
“aberration” could not possibly be due to “parallax,” it may be stated
that aberration shifts the apparent place of a star in one direction,
while parallax shifts it in the opposite direction.
From photometric comparisons, the Rev. John Mitchell, in the year 1767,
concluded that the parallax of Sirius is less than a second of arc; a
result which has been fully confirmed by modern measures. He considered
that stars of the sixth magnitude are probably 20 to 30 times the
distance of Sirius, and judging from their relative brilliancy alone,
this result would also be nearly correct. But recent measures have shown
that some of the fainter stars are actually nearer to us than some of
the brighter, and that the brightness of a star is no criterion of its
distance.
The first stars on which observations seem to have been made with a view
to a determination of their distance seem to have been Aldebaran and
Sirius. From observations made in the years 1792 to 1804 with a vertical
circle and telescope of 3 inches aperture, Piazzi found for Aldebaran an
“absolute” parallax of about 1½ seconds of arc. O. Struve and Shdanow,
in 1857, using a refractor of 15 inches aperture, found a “relative”
parallax of about half a second. This was further reduced by Hall with
the 26-inch refractor of the Washington Observatory to about one-tenth
of a second, and Elkin, with a heliometer of 6 inches aperture, finds a
relative parallax of 0″·116, or about 30 years’ journey for light For
Sirius, Piazzi found, in 1792–1804, an absolute parallax of four
seconds, but this was certainly much too large. All subsequent observers
find a much smaller parallax, recent measures giving a relative parallax
of 0·370″ by Gill, and 0·407″ by Elkin. In the years 1802–1804, Piazzi
and Cacciatori found an absolute parallax of 1′·31 for the Pole Star;
but this has been much reduced by other observers. Pritchard, by means
of photography, found a relative parallax of only 0·073″, which agrees
closely with some other previous results, and indicates a “light
journey” of about 44 years!
For the bright star Procyon, Piazzi found a parallax of about three
seconds, but this is also much too large, a recent determination by
Elkin giving 0·266″, a figure in fair agreement with results found by
Auwers and Wagner. For the bright star Vega, Calandrelli, in the years
1805–6, found an absolute parallax of nearly four seconds, but this has
also been much reduced by modern measures; Elkin, from observations in
the years 1887–88, finding a relative parallax of only 0·034″. Brinkley
found a parallax of over one second for Arcturus, but Elkin’s result is
only 0·018″. If this minute parallax can be relied on, Arcturus must be
a sun of vast size.
Owing to the large “proper motion” of the star known as 61 Cygni, its
comparative proximity to the earth was suspected, and in 1812, Arago and
Mathieu found, from measures made with a repeating circle, a parallax of
over half a second. Various measures of its parallax have since been
made, ranging from about 0·27″ to 0·566″. Sir Robert Ball, at Dunsink,
Ireland, found 0·468″, and Pritchard, by means of photography with a
13-inch reflector, found 0·437″. We may, therefore, safely assume that
the parallax of 61 Cygni is about 0·45″. This implies a distance of
458,366 times the sun’s distance from the earth, or about 42 billions of
miles, and a “light journey” of about 7¼ years.
It is usually stated that 61 Cygni is the nearest star to the earth in
the Northern Hemisphere, but for the star known as Lalande 21,185,
Winnecke found 0·511″, and afterwards 0·501″. This has, however, been
reduced by Kapteyn (1885–1887) to 0·434″; and recently a parallax of
0·465″ has been found by the photographic method for the binary star,
Eta Cassiopeiæ. 61 Cygni is a wide double star, but it seems doubtful
whether the components are physically connected, although several orbits
have been provisionally computed.
Nearer to us than 61 Cygni is the bright southern star Alpha Centauri,
which, so far as is known at present, is the nearest of all the fixed
stars to the earth. The first attempt to find its distance was made by
Henderson in the years 1832–33, using a mural circle of 4 inches
aperture and a transit of 5 inches. He found an “absolute” parallax of
about one second of arc, which subsequent measures have shown to be
rather too large. Measures in recent years range from 0·512″ to 0·976″,
but probably the most reliable are those made with a heliometer of 4½
inches aperture by Dr. Gill (1881–82), who found a “relative” parallax
of 0·76″, and by Dr. Elkin, using the same instrument, 0·671″. Gill’s
result would place the star at a distance of 271,400 times the sun’s
distance from the earth, or about 25 billions of miles, a distance which
light, with its great velocity of 186,300 miles a second, would take
over 4¼ years to traverse.
It will be understood that the parallaxes found for even the nearest
fixed stars are so small that their exact determination taxes the powers
of the most perfect instruments and the skill of the most experienced
observers. One thing, however, seems certain, that the brightest stars
are not necessarily the nearest, and that comparatively faint stars may
be actually nearer to the earth than some of the brightest gems which
deck our midnight sky. Indeed, from a discussion of the observed
parallaxes and “proper motions” of 11 stars, Gylden finds a mean
parallax of only 0·083″ for stars of the first magnitude. This agrees
closely with the value 0·089″ found by Dr. Elkin.
In old times the stars were supposed to be absolutely fixed in the
celestial vault, that is to say, that their relative positions did not
change. This was a very natural conclusion, for before the invention of
the telescope it would have been impossible to detect any “proper
motion”—as it is called—by naked eye observations. Hence the term “fixed
stars,” used to distinguish the stars from the planets, which are always
shifting their positions in the heavens. The existence of proper motion,
in some at least of the stars, seems to have been discovered by Halley,
who found from his observations in 1715 that the bright stars, Sirius,
Arcturus, and Aldebaran, had apparently shifted their positions since
the date of the earliest observations. This discovery was confirmed by
James Cassini in 1738. He found that Arcturus had apparently moved
through some five minutes of arc in 152 years, or about two seconds a
year, a result which agrees fairly well with more exact modern measures.
This interesting discovery of stellar motion has been fully confirmed by
modern observations, and we now know that, far from the stars being
“fixed,” most of them have an apparent motion on the celestial vault.
These motions are, however, very slow, and can only be detected by
accurate measurements and a careful comparison of their positions after
the lapse of a number of years. The largest proper motion hitherto
detected is that of a star known as 1830 of Groombridge’s catalogue, a
small star of about 6½ magnitude, which lies in the constellation Ursa
Major. This star has an apparent motion of seven seconds per annum,
which, though relatively large, is of course absolutely small, as the
observed motion would only suffice to carry it through a space equal to
the moon’s apparent diameter in about 266 years. Assuming a parallax of
about one-sixth of a second found by Kapteyn, this apparent motion would
indicate a real motion of about 128 miles a second at right angles to
the line of sight. As, however, there may be also motion _in_ the line
of sight, the above velocity would be a minimum—if the parallax can be
relied upon—and the actual motion may be considerably more. From its
rapidity, 1830 Groombridge has been called by Prof. Newcomb “the runaway
star.”
Next in order of rapidity of motion comes the southern star known as
Lacaille 9352, which lies in the constellation Piscis Australis, a
little south of Fomalhaut. This seventh magnitude star has an apparent
motion of 6·9 seconds, which, with a parallax of 0·285″ found by Gill,
indicates a velocity of 71 miles per second. Next comes 61 Cygni, with a
velocity of 30 miles, and Epsilon Indi—another southern star—with a
velocity of nearly 68 miles a second. These velocities are, however,
exceeded by other stars if the measured parallaxes are correct. Thus the
star Mu Cassiopeiæ, with a proper motion of 3·7 seconds, has, according
to Pritchard’s photographic measures, a parallax of only 0·036″, which
would indicate a velocity of no less than 302 miles a second! and the
small parallax found by Elkin for Arcturus would imply the startling
velocity of 376 miles a second!
It is a remarkable fact that the eight stars with the largest proper
motions are all below the fourth magnitude in brightness, and as a large
proper motion probably indicates proximity to the earth, the conclusion
seems evident that the brightest stars are not as a rule the nearest. Of
twenty-five stars, with proper motions greater than two seconds of arc,
there are only two—Arcturus and Alpha Centauri—whose magnitude exceeds
the third. Indeed, more than half the stars with motions greater than
one second are invisible to the naked eye!
Many stars have proper motions of less than a second of arc per annum.
Very small proper motions have also been detected, which only reveal
themselves after the lapse of a great number of years, and it seems
probable that there are no really “fixed stars” in the heavens. For
stars of the sixth magnitude, M. Ludwig Struve finds an average motion
of only eight seconds in a hundred years, or about one-twelfth of a
second per annum. If we assume that stars of the sixth magnitude are, on
the average, of the same size and brightness as stars of the first
magnitude, their distance from the earth would be ten times greater.
Consequently, stars of the first magnitude should have an average proper
motion of about eighty seconds in one hundred years. This, however, is
not the case. The twenty brightest stars show an average motion of only
sixty seconds in a hundred years. And the motion of stars of the second
magnitude is relatively still slower. Instead of an average motion of
fifty seconds in a hundred years—which they should have if the
brightness were inversely proportional to the distance—it has been found
that twenty-two stars of the second magnitude show an average motion of
only seventeen seconds. This result seems to show that the brighter
stars are not so near us as their brilliancy would lead us to suppose, a
conclusion which has been already proved by actual measures of their
distance.
From a consideration of the results found for stellar parallax, Mr.
Thomas Lewis, F.R.A.S., of the Greenwich Observatory, comes to the
following conclusions[108]:—
“(1) Leaving out a few of the brightest stars, the parallaxes are
constant down to 2·70 magnitude.
“(2) After 2·70 mag. is reached, the parallaxes are doubled, and remain
practically constant to 8·40 mag.
“(3) Up to the 3rd mag. the velocities are very small, averaging about 9
miles per second, while after the 3rd mag. the velocity is 38 miles per
second.
“Hence we may fairly deduce—
“(1) That there are a few stars (about 8) of exceptional brilliancy in
our immediate neighbourhood, and scattered about amongst these a number
of small stars (at present about 40 are known).
“(2) Stars of mag. 1·0 to 3·0 are, as a class, far outside this inner
space, and have very small velocities.
“(3) The small stars here dealt with have apparently large velocities
across the line of sight.
“These results show that the generally received idea that parallaxes are
to be sought for in stars with large proper motion is correct, and we
may add that this holds good, no matter what may be the star’s
magnitude.”
The “proper motion” of a star only indicates its motion at right angles
to the line of sight—that is, its motion on the surface of the celestial
vault—and gives us no information as to whether the star is approaching
to or receding from the earth. This motion “in the line of sight” cannot
be detected by micrometrical measures with an ordinary telescope, and
would probably have remained for ever unknown had the spectroscope not
been invented. Dr. Huggins was the first to show that motions in the
line of sight could be determined by measuring the displacement of the
spectral lines caused by the approach or recession of the source of
light, the lines being slightly shifted towards the blue end of the
spectrum when the star is approaching the earth, and towards the red end
when it is receding from us. The effect would, of course, be exactly the
same if the star were at rest and the earth in motion. By carefully
measuring this observed displacement of the spectral lines, the velocity
in the line of sight can be easily computed. Dr. Huggins’ observations
were fully confirmed by Dr. Vogel.
The earlier determinations of motion in the line of sight were made by
eye measurements with a micrometer, and owing to the difficulty and
delicacy of these measures, the results were very discordant. The method
has recently been much improved by photographing the spectra and
measuring the positions of the lines on the photograph. Both methods
agree in showing that the following stars, among others, are certainly
_approaching_ the earth: Arcturus, Vega, Procyon, Pollux, Altair, Spica,
Alpha Cephei, Alpha Persei, Alpha Arietis, 61 Cygni, and the Pole Star;
and the following are certainly _receding_: Capella, Rigel, Betelgeuse,
Aldebaran, and Regulus.
Measures of photographic stellar spectra have yielded much more accurate
results than the old method. Some of the velocities found in this way by
Dr. Vogel—who has given especial attention to this subject—are very
considerable. For the bright star Rigel he finds a velocity of recession
of about 39 miles a second, for Aldebaran 30 miles, and for Capella 15
miles. He finds that the Pole Star is approaching the earth at the rate
of 16 miles a second, and Procyon about 7 miles.
Dr. Bélopolsky has recently investigated the _absolute_ velocity in
space of the brighter component of 61 Cygni—that is, the motion across
the line of sight combined with the motion _in_ the line of sight.
Assuming a parallax of half a second and a proper motion of 5·2 seconds,
he finds that the motion across the line of sight, corrected for the
sun’s motion in space, is about 22½ miles per second. The motion _in_
the line of sight, also corrected for the sun’s motion, he finds, from
photographs taken at Pulkova, to be about 27 miles a second towards the
earth. Combining these motions, he finds the absolute velocity of the
star in space to be about 35 miles a second, or nearly double the
velocity of the earth in its orbit
This method of measuring velocities in the line of sight has also been
applied to the nebulæ. Mr. Keeler has observed and measured a
displacement of the line known as the chief nebular line in several
planetary nebulæ, and finds considerable motion in the line of sight.
For example, in the nebula numbered 6790 in the “New General Catalogue,”
he finds a motion of recession of about 38 miles a second. Some of these
motions may possibly be due, in part at least, to the sun’s motion in
space, carrying the earth with it, a motion which will now be
considered. The method has also led to the discovery of the so-called
“spectroscopic binary stars,” a most interesting class of objects, which
will be considered in the next chapter.
The proper motions of the stars long since suggested the idea that
possibly the observed motion may be—to some extent, at least—merely
apparent, and due to the real motion of the sun and solar system through
space. The first investigation of this interesting question was made by
Sir William Herschel in 1783, and he came to the conclusion that the sun
is moving towards a point near Lambda Herculis, a result not differing
widely from modern determinations. The reality of Herschel’s result has
been fully confirmed by subsequent investigations, and Argelander placed
it beyond doubt by a comparison of the positions of a large number of
stars determined at Abo with those found by Bradley in 1752. The
accuracy of Argelander’s result was confirmed by Otto Struve. According
to the elder Struve, the results arrived at by Argelander, O. Struve,
and Peters, is to place the point towards which the sun is moving,
between the stars Pi and Mu Herculis, “at a quarter of the apparent
distance of these stars from Pi Herculis,” and they estimated the annual
motion at about 33½ million miles geographical. The general accuracy of
this conclusion has been verified by modern researches, although the
results found by different astronomers vary to some extent. The
accompanying diagram shows some of the different positions found by
various computers. The later determinations seem to place the “apex of
the solar motion,” as it is termed, not far from the bright star Vega,
or further to the east than Herschel placed it. The velocity of the
sun’s motion in space has not been so well determined as its direction.
L. Struve’s computations would indicate a velocity of about 14 miles a
second; but other results give a much smaller velocity.
[Illustration:
FIG. 3.—_Diagram showing “Solar Apex,” and the different Positions
found by various Computers._
(From “Visible Universe.”)
]
From a recent investigation of the nature of the sun’s motion in space
by Mr. G. C. Bompas,[109] he considers that the various positions of the
sun’s “apex” show a tendency to a drift along the edge of the Milky Way,
and that this drift “seems to point to a plane of motion of the sun
nearly coinciding with the plane of the Milky Way, or, perhaps, more
nearly with the plane of that great circle of bright stars first
described by Sir Wm. Herschel as inclined about 20° to the galaxy, and
which passes through Lyra, in or near which constellation the solar apex
lies,” and he concludes, from the motion of the nearer stars, “that the
sun moves in a retrograde orbit from east to west, and in a plane
inclined a few degrees to that of the Milky Way.” With reference to this
very interesting conclusion, which may, perhaps, be confirmed by further
observations, Mr. Bompas quotes the following passage from “The Visible
Universe,” p. 197, by the present writer:—“With reference to a possible
motion of the stars in some general system, M. Rancken has found, from
an examination of 106 stars, a tendency to drift along the course of the
Milky Way from Aquila towards Cygnus and Cassiopeia, and past Capella
through Orion to Argo. The _larger_ motions, shown in Proctor’s map of
‘proper motions,’ exhibit this tendency in a marked degree between
Cygnus and Capella, and less clearly on the Sirius, but the smaller
motions not so well,” and Mr. Bompas points out that this apparent drift
of the stars in the Milky Way, from west to east, “is just such as would
be occasioned by a real motion of the sun in that plane, in a contrary
direction from east to west.”
CHAPTER IV.
BINARY STARS.
Double and multiple stars may be either optical or real. Optical double
stars are those in which the component stars are merely apparently close
together, owing to their being seen in nearly the same direction in
space. Two stars may _seem_ to be close together, while, in reality, one
of them may be placed at an immense distance behind the other. Just as
two lighthouses at sea may, on a dark night, appear close together when
viewed from a certain point, whereas they may be really miles apart. In
the case of double stars it is, of course, always difficult to determine
whether the apparent closeness of the stars is real or merely optical.
But when, from a long series of observations of their relative position,
we find that one is apparently moving round the other, we know that the
stars must be comparatively close, and linked together by some physical
bond of union. These most interesting objects are known to astronomers
as binary or revolving double stars. The probable existence of such
objects was predicted from abstract reasoning by Mitchell in the
eighteenth century; but the discovery of their actual existence was made
by Sir William Herschel, while engaged on an attempt to determine the
distance of some of the double stars from the earth. “Instead of
finding, as he expected, that annual fluctuation to and fro of one
component of a double star with respect to the other—that alternate
increase and decrease of their distance and angle of position, which the
parallax of the earth’s annual motion would produce—he observed, in many
cases, a regular progressive change; in some cases bearing chiefly on
their distance, in others on their position, and advancing steadily in
one direction, so as clearly to indicate a real motion of the stars
themselves,” and measurements made during the subsequent 25 years fully
proved the truth of the illustrious astronomer’s discovery. It was found
that in many double stars an orbital motion round each other was evident
after a number of years of careful observation of their relative
positions. Unlike the planetary orbits, which are nearly circular, at
least those of the larger planets of the solar system, it was found that
the orbits of these double stars differ, in many cases, widely from the
circular form, in some cases, indeed, approaching in shape more the
orbit of a comet than a planet.
The binary stars are among the most interesting objects in the heavens.
The number now known probably amounts to nearly one thousand. In most of
them, however, the motion is very slow, and in only about seventy cases
has the change of position, since their discovery, been sufficient to
enable an orbit to be computed. In most cases the plane of the real
orbit, or ellipse, described by the companion round the principal star,
is inclined to the line of sight. We therefore see the orbit
foreshortened into a more elongated ellipse.
The relation of the apparent ellipse—or the ellipse we see described by
one star round the other—to the real ellipse will be easily understood
by the following illustration. Suppose a cylinder or rod of an
elliptical, not circular, section to be cut across obliquely to its
axis. This oblique section will represent the _real_ orbit of a binary
star, and the section at right angles to the axis, the _apparent_ orbit.
The angle between these two sections will represent the inclination of
the real orbit to the plane of projection, or background of the sky. In
the apparent orbit, the primary star, which is assumed to be situated in
one of the foci of the real ellipse, does not lie in the focus of the
apparent ellipse, and from its observed position in this latter ellipse
we can deduce, mathematically, the particular angle at which the oblique
section must be made to agree with the observed place of the primary
star, and other details respecting the real ellipse.
Savary, in 1830, was the first astronomer who attempted to compute the
orbit of a binary star, namely, the star Xi Ursæ Majoris. This
remarkable pair was discovered by Sir William Herschel in 1780, and as
the period of revolution is about 61 years, a considerable portion of
the ellipse had been described in 1830, when it was attacked by Savary.
Since that year, orbits have been computed for a number of binary stars
by several computers, among whom may be mentioned Sir R. Ball, Behrmann,
Casey, Celoria, Doberck, Dunér, Elkin, Fritsche, Glasenapp, Sir J.
Herschel, Hind, Jacob, Mädler, Mann, Schur, See, Thiele, Villarceau, and
the present writer. The computation of a double star orbit is a matter
of considerable trouble and difficulty, and cannot be described here. An
account of the principal results arrived at by astronomers in this
interesting branch of sidereal astronomy may, however, prove of interest
to the general reader.
We will first consider the binary stars with short periods of
revolution, which are, of course, the most interesting, and those whose
orbits can be computed with greater accuracy than binaries having
periods of considerable length. The binary star with the shortest period
known at present seems to be the fourth magnitude star Kappa Pegasi. It
was discovered as a wide double star by Sir William Herschel in 1786,
the companion star being of the ninth magnitude. In August, 1880, Mr.
Burnham, the famous American double star observer, examining the star
with the 18½ inch refractor of the Dearborn Observatory, found the
brighter star to be a very close double, with a distance between the
components of only a quarter of a second of arc. A few years’
observations showed that this pair were in rapid motion round each
other, and from measures up to the year 1892, Burnham finds a period of
11·37 years. A later determination by Dr. See makes the period 11·42
years, so that we may conclude that the orbit is now pretty accurately
determined. The plane of the orbit is highly inclined to the line of
sight. Dr. See makes the inclination 81°.
Another binary star, with a period of about the same length, is Delta
Equulei, which was discovered to be a close double by Otto Struve in
1851. As in the case of Kappa Pegasi, the orbit is highly inclined to
the line of sight. In the year 1887, Wrublewsky, the Russian computer,
found a period of about 11½ years, with an orbit nearly circular. A new
orbit was published in 1895 by Dr. See, who finds a period of 11·45
years, and an orbit agreeing fairly well with that of Wrublewsky, the
orbit differing little from the circular form, and inclined to the line
of sight at the high angle of 79 degrees. Burnham found only a “slight
elongation” in the star with the great 36-inch telescope of the Lick
Observatory in July, 1889. The distance between the components does not
at any time exceed half a second of arc, so that it is always beyond the
reach of all but the largest telescopes.
Next in order of shortness of period comes the southern binary star Zeta
Sagittarii, for which an orbit was first computed in the year 1886 by
the present writer, who found a period of 18·69 years. The orbit was
re-computed in 1893, with the aid of recent measures by Mr. J. W.
Froley, who finds a period of 17·71 years. The orbit of this star will,
I think, require still further revision, but the period of about 18
years is probably not far from the truth.
Another remarkably rapid binary star is 85 Pegasi, for which Schaeberle
computed a period of 22·3 years, but a later orbit by Prof. Glasenapp
makes the period 17½ years, and Burnham thinks it will certainly be less
than 20 years. Dr. See, however, finds a period of 24 years. The primary
star is about the sixth magnitude, and the companion only the eleventh,
a difference of five magnitudes, which implies that the larger star is
one hundred times brighter than the companion.
Next in order of rapidity of motion we have the southern binary star 9
Argûs. For this pair, Burnham finds a period of 23·3 years, and Dr. See
22 years, the other elements of the orbit being also in close agreement.
In this case also the orbit plane is highly inclined to the line of
sight.
The star 42 Comæ Berenices has a period of about 25¾ years, according to
Otto Struve. The orbit is remarkable from the fact that its plane passes
through or nearly through the earth, and is, therefore, projected into a
straight line, the companion star oscillating backwards and forwards on
each side of its primary. I find that the plane of the orbit is at right
angles to the general plane of the Milky Way.
The star Beta Delphini—the most southern of the four stars in the
“Dolphin’s Rhomb”—is also a fast-moving binary, discovered by Burnham in
1873, for which periods have been computed of 22·97 years by Glasenapp,
26·07 years by Dubjago, 27·66 years by Dr. See, and 30·91 years by the
present writer. Burnham thinks the period will prove to be about 28
years. The spectrum of the light of Beta Delphini is similar to that of
our sun, so that the two bodies should be comparable in intrinsic
brilliancy. From my orbit of the pair, the “hypothetical parallax” is
0·052″—that is, this is the parallax the star would have on the
supposition that the combined mass of its components is equal to the
mass of the sun. Now, assuming the value of the sun’s stellar magnitude
which I have recently computed (_Knowledge_, June, 1895)—namely, 27·15—I
find that the sun, if placed at the distance indicated for Beta
Delphini, would be reduced to a star of 5·84 magnitude. As the star was
measured 3·74 at Harvard, we have a difference of 2·1 magnitude,
denoting that the binary—if of the same mass as the sun—must be about
seven times brighter. As the spectrum is of the same type, this seems
improbable, and we must conclude that the star’s parallax is more than
0·052″.
Another remarkable binary star with a comparatively short period is Zeta
Herculis. This pair have now performed three complete revolutions since
their discovery in 1782 by Sir William Herschel. Several orbits have
been computed, but Dr. See’s period of 35 years is probably the best The
companion is now not far from its maximum distance (1½ seconds) from the
primary star, and is within the reach of moderate-sized telescopes. The
companion is, however, rather faint, being only 6½ magnitude, while the
primary star is of the third. When at their nearest, some observers have
spoken of an “occultation” of one star by the other, but no real
occultation ever takes place, the components never approaching within
half a second of arc. The companion merely disappears owing to its
faintness in telescopes of moderate power. An occultation of one
component of a binary star by the other cannot take place except—as in
the case of 42 Comæ—when the plane of the orbit passes through the
earth.
[Illustration:
FIG. 4.—_Apparent Orbit of Zeta Herculis._ (From “Worlds of Space.”)
]
In the case of the binary star, Eta Coronæ Borealis, it was, some forty
years ago, uncertain whether its period was 43 or 66 years, but now that
two complete revolutions have been performed since its discovery by Sir
William Herschel in 1781, the question has been finally decided in
favour of the shorter period. Numerous orbits have been computed, but
these by Dr. Doberck and Dr. Dunér are probably the best. Those give a
period of about 41½ years. The components are nearly equal in
brightness, but at their present distance are not within the reach of
small telescopes.
The brilliant star Sirius is also an interesting binary star. The
companion, which is relatively very faint—about tenth magnitude,—was
discovered by Alvan Clark in 1862. The existence of some such disturbing
body was previously suspected by astronomers, owing to observed
irregularities in the proper motion of Sirius. Several orbits, giving
periods of about 50 years, have been computed. Some measures in recent
years, however, seemed to show that this period was somewhat too short,
but a period of about 58½ years, computed by the present writer in 1889,
will probably prove too long. Some few years ago, Burnham found the
companion an easy object with the 36-inch refractor of the Lick
Observatory, but towards the end of the year 1890 it passed beyond the
power of even this giant telescope. It will probably, however, emerge
very soon now from the rays of its brilliant primary.[110] Burnham finds
a period of about 52 years, but the German astronomer, Auwers, who has
carefully investigated the observed irregularities in the proper motion
of Sirius, adheres to a period of about 49½ years. The great brilliancy
of Sirius, the brightest star in the heavens, naturally suggests a sun
of great size. Recent investigations, however, do not favour this idea.
Assuming a parallax of 0″·39 (about a mean of the results found by Elkin
and Gill), Auwers finds the mass of the system to be about three times
the mass of the sun, the mass of the companion being about equal to the
sun’s mass. Placed at the distance of Sirius, the sun would, I find, be
reduced to a star of about 1½ magnitude. As Sirius is about 1 magnitude
brighter than the zero magnitude—that is, about 2 magnitudes brighter
than a standard star of the first magnitude—it follows that it is about
2½ magnitudes, or about ten times brighter than the sun would be in the
same position. Its spectrum is, however, of the first type, and the star
is therefore not comparable with the sun in brilliancy. The above result
would indicate that stars of the first or Sirian type are intrinsically
brighter than our sun.
Sirius is about 11 magnitudes brighter than its faint companion. This
makes the light of Sirius about 25,000 times the light of the small
star. If, therefore, the two bodies were of the same intrinsic
brilliancy, their diameters would be in the ratio of 158·5 to 1, and if
of the same density, the mass of Sirius would be nearly five million
times the mass of the companion! But, according to Auwers’ calculations,
the companion’s mass is about one-half that of its primary. The two
bodies must, therefore, be differently constituted, and, indeed, the
companion must be nearly a dark body. It has been suggested that the
companion may possibly shine by reflected light from Sirius; but this I
have shown elsewhere to be quite impossible.[111] Even with a diameter
equal to that of the sun, I find that with reflected light only it would
be quite invisible in all parts of its orbit, even with the great Lick
telescope. It must, therefore, shine with inherent light of its own, and
it seems probable that it is a large body, cooling down and approaching
the complete extinction of its light. If Sirius has any planets
revolving round it—like those of our solar system—they must for ever
remain invisible in our largest telescopes. This remark, of course,
applies to all the fixed stars, single and double. They may possibly
have attendant families of planets, like our sun, but if so, the fact
can never be ascertained by direct observation. I find that the plane of
the orbit of Sirius is at right angles to the general plane of the Milky
Way.
[Illustration:
FIG. 5.—_Apparent Orbit of the Companion of Sirius._
(From “Old and New Astronomy.”)
]
The star Zeta Cancri is a well-known triple star, the close pair
revolving in a period of about 60 years. Nearly two revolutions have now
been completed since its discovery by Sir William Herschel in 1781. All
three stars probably form a connected system, but the motion of the
third star round the binary pair is very slow and irregular. The motion
of this interesting system has recently been investigated by Professor
Seeliger, and he comes to the conclusion that, to make the observations
agree with calculation, it is necessary to assume that the third star is
in reality a very close double, the components of which revolve round
their centre of gravity in about 17½ years, and both round the known
binary pair. If this be so, we have here a remarkable quadruple pair;
but it must be added that all efforts with large telescopes to see the
companion star double have failed, and that the existence of the fourth
star rests only on theory. Burnham, in 1889, using a power of 1500,
failed to see any other component.
Another interesting binary star is Xi Ursæ Majoris. As already stated,
this was the first pair for which an orbit was computed. More than a
complete revolution has now been performed since its discovery by Sir
William Herschel in 1780. The period has, therefore, been well
determined, and seems to be about 60 years. Although the components are
not near their maximum distance at present, they are still within the
reach of moderate telescopes, the distance being about 1¾ seconds, and
the magnitudes of the components, not very unequal, about 4 and 5.
The bright southern star, Alpha Centauri, the nearest of all the fixed
stars to the earth, so far as is known at present, is also a remarkable
binary star. It seems to have been first noticed as a double star by
Richaud in 1690. Several orbits have been computed, ranging from about
75 to 88½ years, but recent calculations by Mr. A. W. Roberts and Dr.
See make the period about 81 years, which agrees closely with Dr.
Elkin’s period of 80⅓ years. Combining Dr. Gill’s parallax of 0″·76 with
Elkin’s elements, I find the sum of the masses nearly twice the mass of
our sun, and the mean distance between the components about 23 times the
earth’s distance from the sun, or somewhat greater than the distance
between the sun and Uranus. Dr. Doberck finds a period of about 79
years, and assuming a parallax of 0″·75, he finds the mean distance
between the components 24·6 times the earth’s distance from the sun; and
he points out that if we suppose that their diameter does not differ
much from that of our sun, each component “would appear from the other
as a mere star to unaided vision, the distance being too great to show a
disc.”[112] From a recent investigation of the proper motion and
position of Alpha Centauri, Mr. A. W. Roberts finds that the masses of
the components are nearly equal, and the combined mass equal to twice
the mass of our sun, a conclusion in close agreement with the result
found above from the orbit. According to Dr. Gill, the difference in
brightness of the two components is 1·25 magnitude, and Professor Bailey
makes their photometric magnitudes 0·50 and 1·75. As this difference
would make the brighter component over three times brighter than the
companion, it follows that its surface must be much brighter, and Mr.
Roberts concludes that the companion has proceeded “some distance on the
down track from a sun to an ordinary planet.” Assuming my value of the
sun’s stellar magnitude (about 27), I find that the sun, if placed at
the distance of Alpha Centauri, would appear of about the same
brightness as the star does to us. As, according to Professor Pickering,
the spectrum of Alpha Centauri is of the second or solar type, it would
seem that in mass, brightness, and physical condition, the star closely
resembles our sun.
We next come to another very interesting binary star, known to
astronomers as 70 Ophiuchi. It is a very fine double star, the
magnitudes of the components being about 4 and 6, and the colours yellow
and orange. More than a complete revolution has now been described by
the components since its discovery by Sir William Herschel in 1779.
Numerous orbits have been computed with periods ranging from 73¾ to 98
years. An orbit computed by the present writer, in 1888, gave a period
of 87·84 years, and this was confirmed in 1894 by Burnham, who found a
period of 87·85 years. A subsequent investigation by Schur gives a
period of 88·356 years. My orbit, combined with Krüger’s parallax of
0″·162, give for the combined mass of the components 2·777 times the
mass of the sun, and the distance between them 27·777 times the earth’s
distance from the sun, or somewhat less than the distance of Neptune
from the sun. Schur has, however, recently found a parallax of 0″·286,
which would reduce the mass of the system, and also the distance between
the components. Recent observations show that the companion is now in
advance of the theoretical position indicated by Schur’s orbit, and Dr.
See thinks that the observed irregularities in the orbital motion of the
pair indicate the existence of a third body, and that either the primary
star or the companion, probably the latter, is a very close binary star.
Careful search, however, for a third body, made with large telescopes,
have failed to reveal its existence, and so the matter remains in
suspense. Placed at the distance indicated by Krüger’s parallax, I find
that our sun would be reduced to a star of about magnitude 3½, which
shows that the sun and star are of about equal brightness. The spectrum
is of the solar type, according to Vogel. I find that the plane of the
orbit is at right angles to the plane of the Milky Way.
The star Gamma, in Corona Borealis, is a close and difficult binary
star. Dr. Doberck finds a period of 95½ years, and Celoria about 85¼. As
in the case of 42 Comæ, the plane of the orbit nearly passes through the
earth, and the apparent orbit is, consequently, nearly a straight line.
I find that the plane of the orbit is at right angles to the plane of
the Milky Way.
The star Xi Scorpii is a remarkable triple star, like Zeta Cancri, the
magnitudes of the components being about 4½, 5, and 7½. The components
of the close pair have described a complete revolution since their
discovery by Sir William Herschel in 1780. Dr. Doberck finds a period of
about 96 years, and Schorr 105 years. The real orbit is nearly circular,
but owing to its high inclination, about 70°, the apparent orbit is a
very elongated ellipse. All three stars have a common proper motion
through space, and, probably, form one system, but the motion of the
third star is very slow, and its period of revolution must be several
hundred years.
[Illustration:
APPARENT ORBIT OF 70 OPHIUCHI, COMPUTED BY J. E. GORE (1888).
(_Showing positions of companion star in
different years._)
(From “The Scenery of the Heavens.”)
]
The star ο^2, or 40 Eridani, is another interesting object. It is a star
of about 4½ magnitude, with a distant ninth magnitude companion, which
is a double and binary star. It is sometimes stated that the bright star
is the binary, but this is quite incorrect; the large star is single—at
least, as far as is known at present. An orbit for the binary pair was
computed, in 1886, by the present writer, who found a period of 139
years; but Burnham, using later observations, finds a period of 180
years. A physical connexion may possibly exist between the binary pair
and the bright star, as both have the same common motion through space,
but the angular motion, if any, is very slow. Professor Asaph Hall found
a parallax of about one-fifth of a second of arc, and this, combined
with Burnham’s orbit, gives the combined mass of the binary pair about
two-thirds of the sun’s mass, a result which seems remarkable, for the
sun, placed at the distance indicated by Hall’s parallax would, I find,
shine as a star of about the third magnitude, or considerably brighter
than the principal star of 40 Eridani. Owing to the faintness of the
binary pair, the nature of its spectrum has not been determined.
Computed by a well-known formula, its “relative brightness”—that is, its
brightness compared with that of other binaries—is very small.
A very famous binary star is that known to astronomers as Gamma
Virginis. Its history is a very interesting one. It lies close to the
celestial equator, about one degree to the south and about fifteen
degrees to the north-west of the bright star Spica (Alpha of the same
constellation), with which it forms the stem of a Y-shaped figure,
formed by the brightest stars of the constellation Virgo, or the Virgin,
Gamma being at the junction of the two upper branches. The brightness of
Gamma Virginis is a little greater than an average star of the third
magnitude. Photometric measures made at Oxford and Harvard Observatories
agree closely, and make its brightness about 2·7 magnitude—that is to
say, rather nearer the third than the second magnitude. Variation of
light has, however, been suspected in one or both components, and this
question of light variation will be considered further on. The Persian
astronomer, Al-Sûfi, in his description of the heavens, written in the
tenth century, rates it of the third magnitude, and describes it as “the
third of the stars of _al-auvâ_, which is a mansion of the moon,” the
first and second stars of this “mansion” being Beta and Eta Virginis,
the fourth star Delta, and the fifth Epsilon, these five stars forming
the two upper branches of the Y-shaped figure above referred to. Gamma
was called _Zawiyah-al-auvâ_, “the corner of the barkers!” perhaps from
its position in the figure, which formed the thirteenth Lunar Mansion of
the old astrologers. It was also called _Porrima_ and _Postvarta_ in the
old calendars. These ancient names of the stars are curious, and their
origin doubtful.
The fact that Gamma Virginis really consists of two stars very close
together seems to have been discovered by the famous astronomer,
Bradley, in 1718. He recorded the position of the components by stating
that the line joining them was then exactly parallel to a line joining
Alpha and Delta of the same constellation. This was, of course, only a
rough method of measurement, and the position thus found by Bradley
being probably more or less erroneous, has given much trouble to
computers of the orbit described by the component stars round each
other, or, rather, round their common centre of gravity. Bradley does
not give the apparent distance between the component stars; but we may
conclude from the orbit, which is now well determined, that they were
then at nearly their greatest possible distance apart. It is curious
that between Bradley’s time and 1794, the star was on several occasions
occulted by the moon; but none of the observers refer to its duplicity.
It was again measured by Cassini in 1720, by Tobias Mayer in 1756, and
by Sir William Herschel in 1780. These measures showed that the distance
between the components was steadily diminishing, and that the position
angle of the two stars was also decreasing. This decrease in the
position angle—measured from the north round by the east, south, and
west, from 0 to 360°—shows that the apparent orbital motion is what is
called retrograde, or in the direction of the hands of a clock, direct
or “planetary motion” being in the opposite direction. The star was
again measured by Sir John Herschel and South in the years 1822–38, by
Struve in the same years, and by Dawes and other observers from 1831 to
the present time. The recorded measures are very numerous, and have
enabled computers to determine the orbit with considerable accuracy. The
rapid decrease in the apparent distance from 1780–1834 indicated that
the apparent orbit is very elongated, and that possibly the two stars
might “close up” altogether, and appear as a single star even in
telescopes of considerable power. This actually occurred in the year
1836, or, at least, the stars were then so close together that the most
powerful telescopes of that day failed to show Gamma Virginis as
anything but a single star. Of course, it would not have been beyond the
reach of the giant telescopes of our day. From the year 1836 the pair
began to open out again, and at present the distance is again
approaching a maximum. It is now within the reach of small telescopes,
and forms a fine telescopic object with a moderate-sized instrument.
The general character of the orbital motion may be described as
follows:—In 1718, at the time of Bradley’s observation, the companion
star was to the north-west of the primary star; it then gradually moved
towards the west and south, and in 1836, when at its minimum distance,
it was to the south-east. From that date it again turned towards the
north, and at present it is north-west of the primary star, and not far
from the position found by Bradley in 1718.
The first to attempt a calculation of the orbit described by this
remarkable pair of suns was Sir John Herschel, who in the year 1831
found a period of about 513 years. In 1833, he re-calculated the orbit,
and found nearly 629 years. We now know that both these periods are much
too long; but the data then available were insufficient for the
calculation of an accurate orbit. From these results Herschel predicted
that “the latter end of the year 1833, or the beginning of the year
1834, will witness one of the most striking phenomena which sidereal
astronomy has yet afforded, _viz._, the perihelion passage of one star
round another, with the immense angular velocity of between 60° and 70°
per annum, that is to say, of a degree in five days. As the two stars
will then, however, be within little more than half a second of each
other, and as they are both large and nearly equal, none but the very
finest telescopes will have any chance of showing this magnificent
phenomenon. The prospect, however, of witnessing a visible and
measurable change in the state of an object so remote, in a time so
short, may reasonably be expected to call into action the most powerful
instrumental means which can be brought to bear on it.” This prediction
was not verified until the year 1836, when the pair “closed up out of
all telescopic reach,” except at the Dorpat Observatory, where a
magnifying power of 848 still showed an elongation in the telescopic
disc of the star. The orbit found by Sir John Herschel was a tolerably
elongated ellipse, with its longer axis lying north-east and south-west.
This was not quite correct, for we now know that this axis lies
north-west and south-east, and that the apparent orbit is much more
elongated than Sir John Herschel at first supposed. This was soon
recognised by Herschel himself, and he came to the conclusion that he
and other computers had been misled by Bradley’s observation in 1718. He
then rejected this early, and evidently faulty, observation, and using
the measures up to 1845, he found a period of about 182 years, which we
now know to be near the truth. The orbit was also computed by the famous
German astronomer, Mädler, who found periods of 145, 157, and 169 years;
by Hind, 141 years; by Henderson, 143 years; by Jacob, 133½, 157½ and
171 years; by Adams, 174 years; by Flammarion, 175 years; and by Admiral
Smyth, 148 and 178 years. All these periods, we now know, are too small.
Fletcher found 184½ years, and Thiele 185 years. Two orbits were
computed by Dr. Doberck, in recent years, with periods of 180½ and 179½
years; but very recently (1895) the orbit has been re-computed by Dr.
See, and he finds a period of 194 years. A comparison of the observed
and computed positions shows, he thinks, that his elements are the most
exact yet determined for any binary star.
The apparent orbit of the pair is a very elongated ellipse, and as
Admiral Smyth said, “more like a comet’s than a planet’s.” The real
ellipse has a very high eccentricity, nearly 0·9—indeed, the greatest of
all the known binary stars, and not much less than that of Halley’s
comet
As I said above, the variability of the light of one or both components
of Gamma Virginis has been strongly suspected. So far back as 1851 and
1852, O. Struve paid particular attention to this point. His
observations in these years show that sometimes the component stars were
exactly equal in brilliancy, and sometimes the southern star—the one
generally taken as the primary—was from 0·2 to 0·7 magnitude brighter
than the other. There seems to be little doubt that some variation
really takes place in the relative brightness of the pair. This is
clearly indicated by the measures of position angle. For example, in the
year 1886, Professor Hall recorded the position as 154·9, evidently
measuring from the northern star as the brightest of the two; while, in
1887, Schiaparelli gives 334°·2—or about 180° more—thus indicating that
he considered the _southern_ star as the primary, or brighter, of the
pair. Burnham found 153°·4 in 1889, and Dr. See 332°·50 in 1891. This is
also shown by earlier measures, for Otto Struve found the southern star
half a magnitude brighter than the other on April 3, 1852, while on
April 29 of the same year he found them “perfectly equal.” He thought
the variation was about 0·7 of a magnitude, but that the climate of
Poulkova, where he observed, was not suitable for such observations.
This variation is very interesting, and the question should be
thoroughly investigated with a good telescope.
As the distance of Gamma Virginis from the earth has not been
determined, it is not possible to calculate the actual dimensions of the
orbit and the mass of the system. If we assume that the combined mass of
the components is equal to the sun’s mass, I find from Dr. See’s orbit
that the “hypothetical parallax” would be 0·119″, implying a distance of
1,733,319 times the sun’s distance from the earth. If, however, we
suppose that the mass of each of the components is equal to the sun’s
mass, or the mass of the system double that of the sun—perhaps a more
probable supposition—I find that the parallax would be about one-tenth
of a second, denoting a distance of 2,062,650 times the sun’s distance
from the earth. Placed at this last distance, the sun would, I find, be
reduced to a star of about 4½ magnitude, or about 1¾ magnitudes fainter
than Gamma Virginis appears to us. This difference implies that,
supposing each of the component stars of the binary to have a mass equal
to the sun’s mass, their combined light is about five times greater than
the sun would emit if placed at the same distance, and as the components
are nearly equal in brightness, each of them would be 2½ times brighter
than the sun. According to Vogel, the star’s light gives a spectrum of
the first or Sirian type, but according to the Draper “Catalogue of
Stellar Spectra,” the spectrum is of the solar type. If the spectrum is
of the first type, its brilliancy is easily explained; for, as I have
shown elsewhere, the Sirian stars, are intrinsically much brighter in
proportion to their mass than those of the solar type. But if its
spectrum is of the solar type, it is not so easy to explain its
brilliancy. Computing by a well-known formula, I find its relative
brightness is nearly five times greater than that of Xi Ursæ Majoris,
the spectrum of which is of the solar type. If, to account for its
brilliancy, we assume that the star is nearer to the earth than the
parallax assumed above would imply, then the mass of the system must be
less than the mass of our sun. As we have seen above, doubling the
supposed mass increased the distance; so, on the other hand, if we
diminish the distance, we must diminish the mass also. Thus, if we
reduce the distance to one-half, we must reduce the mass to one-eighth
of the sun’s mass. A distance of one-third would give a mass of ¹⁄₂₇th,
and a distance of one-fourth would imply a mass only ¹⁄₆₄th of the sun’s
mass. To reduce the sun to the same brightness as Gamma Virginis, it
should be removed to a distance indicated by a parallax of one-tenth of
a second multiplied by the square root of five, or 0·223″. If, however,
the star’s parallax were so much as this, it is probable that it would
have been detected and measured long ago. In the case of the binary star
Castor, I find from the orbit and a small parallax found by Johnson
(about one-fifth of a second) that its mass is only ¹⁄₁₉th of the sun’s
mass, but in this case the spectrum is of the Sirian type, and stars of
this type are very bright in proportion to their mass. The colours of
the components of Gamma Virginis, which are very similar to those of
Castor—white or pale yellow—would suggest that they may belong to the
same type.
Another interesting binary star is Eta Cassiopeiæ. The components are
about 4 and 7½ magnitude, and the pair have described a considerable
portion of their orbit since its discovery in 1779 by Sir William
Herschel, the distance diminishing from about 11 seconds to 4¾. Periods
ranging from 149 to 222½ years have been found by different computers.
The most recent computation makes it about 196 years. Assuming a
parallax of 0·154″ found by Struve, the mass of the system will be from
5¾ to 10¾ times the mass of the sun, according to the length of the
period we assume. A much larger parallax of 0″·3743 was, however, found
by Schweizer and Socoloff, which would considerably reduce the mass, and
recently a still larger parallax of 0″·465 has been found by
photography, which, with Grüber’s elements of the orbit, would reduce
the mass of the system to ⅙th of that of the sun.
The bright star Gamma Leonis, situated in the well-known “Sickle in
Leo,” is also a binary star, but only a small portion of the orbit has
been described since its discovery by Sir William Herschel in 1782. Dr.
Doberck finds a period of 407 years. It is remarkable for its very high
“relative brightness,” which is curious, as its spectrum is of the solar
type. This pair forms a fine object for a small telescope.
The star known as 12 Lyncis is a triple star, the components being 5, 6,
and 7½ magnitude. The close pair form a binary system, for which an
orbit has been computed by the present writer, who finds a period of
about 486 years. Sir John Herschel predicted in 1823 that the angular
motion of the pair would “bring the three stars into a straight line in
57 years.” This prediction was fulfilled in 1887, when measures by
Tarrant showed that the stars were then exactly in a straight line.
[Illustration:
FIG. 7.—_Triple Stars._
(From “Scenery of the Heavens.”)
]
The bright star Castor is a famous double star, and has been known since
the year 1718, when it was observed by Bradley and Pond. It was also
observed by Maskelyne in 1759, and frequently by Sir William Herschel
from 1799 to 1803. Numerous orbits have been computed, with periods
ranging from 199 years by Mädler, and 1,001 years by Doberck. Wilson
found a period of about 983 years, and Thiele about 997 years, so that
the longest period would seem to be nearest the truth. According to a
somewhat doubtful parallax found by Johnson, the distance of Castor from
the earth is about double that of Sirius. With this distance, and
Doberck’s elements of the orbit, I find that the mass of the system of
Castor is only ¹⁄₁₉th of the sun’s mass, a result which would imply that
the components are masses of glowing gas! The spectrum of Sirius is of
the first, or Sirian, type, another example of the great brilliancy of
stars of this type. Quite recently (1896), Dr. Bélopolsky has found,
with the spectroscope, that the brighter component is a close binary
star with a dark companion, like Algol. The period of revolution is
about 3 days, and the relative orbital velocity about 20¾ miles a
second. Dr. Bélopolsky’s observations show that the system is receding
from the earth at the rate of about 4½ miles per second. Assuming the
bright and dark companion to be of equal mass, and hence the absolute
orbital velocity of each one half the relative velocity found by
Bélopolsky, I find that, if the orbit is circular, the distance between
the components is about 85,400 miles, or slightly less than the sun’s
diameter, and their combined mass about ¹⁄₈₇th of the sun’s mass. This
result would imply a still smaller mass for the whole system of Castor
than that found from the orbit of the two bright components, but tends
strongly to confirm the opinion already expressed, that the components
of this remarkable system are merely masses of glowing gas. Assuming
that all three components are of equal mass, the combined mass of the
system would be ¹⁄₅₈th of the sun’s mass. From this result we can easily
compute the stars’ parallax, which, from Dr. Doberck’s orbit, I find to
be 0″·2873, a quantity which might be measured by the photographic
method.
With reference to the colours of the components of binary stars, the
following relation between colour and relative brightness has been
established[113]:—
(1.) When the magnitudes of the components are equal, or approaching
equality, the colours are generally the same, or similar.
(2.) When the magnitudes of the components differ considerably, there is
also a considerable difference in colour.
A new class of binary stars has been discovered within the last few
years by means of the spectroscope. These have been called
“spectroscopic binaries,” and the brighter component of Castor, referred
to above, is an example of the class. They are supposed to consist of
two component stars, so close together that the highest powers of the
largest telescopes fail to show them as anything but single stars.
Indeed, the velocities indicated by the spectroscope show that they must
be so close that the components must for ever remain invisible by the
most powerful telescopes which could ever be constructed by man. In some
of these remarkable objects, the doubling of the spectral lines
indicates that the components are both bright bodies, but in others, as
in Algol, the lines are merely shifted from their normal position, not
doubled, thus denoting that one of the components is a dark body. In
either case, the motion in the line of sight can be measured by the
spectroscope, and we can, therefore, calculate the actual dimensions of
the system in miles, and thence its mass in terms of the sun’s mass,
although the star’s distance from the earth remains unknown. Judging,
however, from the brightness of the star, and the character of its
spectrum, we can make an estimate of its probable distance from the
earth.
Let us first take the case of Algol. This famous variable star has,
according to the Draper catalogue, a spectrum of the Sirian type. It
may, therefore, be comparable with that brilliant star in intrinsic
brightness and density. Assuming the mass of Sirius at 2·20 times the
mass of the sun, as found by Auwers, and that of the brighter component
of Algol at four-ninths of the sun’s mass, as given by Vogel,[114] I
find that for the _same distance_ Sirius would be about 2·8 times
brighter than Algol. But photometric measures show that Sirius is about
22 times brighter than Algol, from which it follows—since light varies
inversely as the square of the distance—that Algol is 2·77 times further
from the earth. Assuming the parallax of Sirius at 0·39″, this would
give for the parallax of Algol O·14″, or a journey for light of about 23
years. From the dimensions of the system, as given by Vogel—about
3,230,000 miles from centre to centre of the components—this parallax
would give an apparent distance between the components of less than
¹⁄₂₀₀th of a second, a quantity much too small to be visible in our
largest telescopes, or probably in any telescope which man can ever
construct From a consideration of irregularities in the proper motion of
Algol and in the period of its light changes, Dr. Chandler infers the
existence of a third dark body and a parallax of 0·07″. As this is
exactly one-half the parallax found above, it implies a distance just
double of what I have found, and would, of course, indicate that Algol
is intrinsically four times brighter than Sirius. This greater
brilliancy would suggest greater heat, and would agree with its small
density, which, from its diameter, as given by Vogel—1,061,000 miles—I
find to be only one-third of that of water.
Let us now consider the case of Beta Aurigæ, which spectroscopic
observations show to be a close binary star with a period of about four
days, and a distance between the components of about eight millions of
miles. This period and distance imply that the mass of the system is
about five times that of the sun. As in this case the spectral lines are
doubled at regular intervals of two days, and not merely shifted, as in
the case of Algol, we may conclude that both the components are bright
bodies, and we may not be far wrong in supposing that they are of equal
mass, each having 2½ times the mass of the sun. As the spectrum of Beta
Aurigæ is of the same type as Sirius, we may compare it with that star,
as we did in the case of Algol. Assuming the same density and intrinsic
brightness for both Beta Aurigæ and Sirius, I find that Beta Aurigæ
should be about twice as bright as Sirius. Now, according to the Oxford
photometric measures, Sirius is 2·89 magnitudes, or 14·32 times brighter
than Beta Aurigæ. Hence it follows that the distance of Beta Aurigæ
should be about 5½ times greater than the distance of Sirius. Hence,
assuming the parallax of Sirius at 0″·39, that of Beta Aurigæ should be
about 0″·061. From actual measures of the parallax of Beta Aurigæ, made
by the late Prof. Pritchard at Oxford, he found, from two companion
stars, a mean parallax of 0″·062, a result in remarkably close agreement
with that computed above from a consideration of the star’s mass and
light, compared with that of Sirius. As the actual distance between the
components of Beta Aurigæ is equal to the sun’s diameter divided by
11·625, we have the maximum angular separation between the components
equal to 0″·062 divided by 11·625, or about ¹⁄₂₀₀th of a second, or
nearly the same as in the case of Algol.
The bright star Spica has also been found by the spectroscope to be a
close binary star. Vogel finds a period of four days with a distance
between the components of about 6¼ millions of miles, and assuming that
the components have equal mass and are moving in a circular orbit, he
finds the mass of the system about 2·6 times the mass of our sun. This
would give each of the components 1·3 times the mass of the sun, and it
follows that the light of Spica—which gives a spectrum of the Sirian
type—should, for equal distances, exceed that of Sirius about 1·4 times.
Now, the photometric measures at Oxford show that Sirius is 1·91
magnitude, or 5·8 times brighter than Spica. Hence it follows that the
distance of Spica should be 2·85 times the distance of Sirius. This
would make the parallax of Spica about 0″·137. So far as I know, a
measurable parallax has not yet been found for this star. Brioschi, in
1819–20, observing with a vertical circle of four inches aperture, found
a negative parallax, which would imply that its parallax is too small to
be measurable. Still, the above result would seem to indicate that its
parallax might be measurable by the photographic method. The parallax
found above would imply that the maximum distance between the components
of Spica would not exceed ⅒th of a second, a quantity much too small to
be detected by the most powerful telescopes. In addition to its orbital
motion, Vogel finds that Spica is approaching the sun at the rate of
over 9 miles per second.
We now come to Zeta Ursæ Majoris (Mizar), which has also a spectrum of
the Sirian type, and which the spectroscopic measures indicate is a
close binary star with a period of about 104 days, and a combined mass
equal to forty times the mass of the sun. Proceeding as before, we find
that the light of Mizar should be about 8·7 times that of Sirius. But
the photometric measures show that Sirius is about three magnitudes, or
about sixteen times brighter than Mizar. Hence the distance of Mizar
should be nearly twelve times the distance of Sirius. This gives for the
parallax of Mizar about 0″·033. Klinkerfues found a parallax of 0″·0429
to 0″·0477, which does not differ widely from the above result. As the
velocity of the orbital motion shown by the spectroscope indicates a
distance between the components of about 143 millions of miles, or about
the distance of Mars from the sun, it follows that the maximum distance
between the components would be 0″032, multiplied by 1½ or 0″·048, a
quantity beyond the reach of our present telescopes.
The well-known variable star, Delta Cephei, has recently been added to
the list of “spectroscopic binaries.” From observations with the great
30-inch refractor of the Pulkowa Observatory in the summer of 1894, M.
Bélopolsky finds that the star is probably a very close double, the
companion being a nearly, or wholly, dark body, as in the case of Algol,
and the orbit a very eccentric one. The observed variation of light
indicates, however, that there is no eclipse, as occurs in Algol, so
that the fluctuations in the light of Delta Cephei are probably due to
some other cause. The spectrum of the star is of the solar type, so that
in this respect it differs from the other spectroscopic binaries
referred to above. The observations show that the system is approaching
the sun at the rate of about 15 miles a second. Spectroscopic
observations also suggest that the well-known variable star Beta Lyræ
may also consist of two close companions. Further details respecting
these observations will be given in the next chapter.
From a recent investigation of the proper motion of the star Tau
Virginis, Dr. Fritz Cohen thinks it is probably a close binary, the
companion star of which has not yet been detected.
It should be mentioned that in the case of Beta Aurigæ, Spica, Zeta Ursæ
Majoris, and Castor, as there is no variation of light, as in Algol, the
plane of the orbit is probably inclined to the line of sight. This would
have the effect of increasing the computed mass of the system, and thus
diminishing the calculated parallax. As the above calculations have been
made on the assumption that the plane of the orbit passes through the
earth, it follows that the computed parallax is a maximum, and that
these remarkable objects may be really further from the earth than even
the minute parallaxes found above would indicate. As the parallaxes of
the nearest stars, such as Alpha Centauri, 61 Cygni, Sirius, and some
other stars, are considerably greater than those found above, it would
seem that our solar system is not situated in a region of binary stars,
and that these wonderful objects lie beyond our immediate neighbourhood.
It is also remarkable that, with the exception of Delta Cephei, they
have all spectra of the Sirian type, including those Algol variables
whose spectra have been examined.
By the aid of the parallaxes computed above, we can easily calculate the
relative brightness of the sun compared with that of the spectroscopic
binaries. Assuming that the sun is 27 magnitudes brighter than the Zero
magnitude, or 28 magnitudes brighter than a standard star of the first
magnitude, and taking the parallax of Algol as 0″·07, I find that the
sun, placed at the distance indicated by this parallax, would be reduced
to a star of 5·35 magnitude, or about three magnitudes fainter than
Algol, which implies that Algol is about 15½ times brighter than our
sun. In the case of Beta Aurigæ, if the sun were placed at the distance
indicated by the parallax of 0″·061, it would be reduced to a star of
5·65 magnitude, or about 3·7 magnitudes fainter than Beta Aurigæ, which
would imply that Beta Aurigæ is about thirty times brighter than the
sun. In the case of Spica we have the sun reduced to a star of about the
fourth magnitude, or about three magnitudes fainter than Spica,
indicating that Spica is, like Algol, about 15½ times brighter than the
sun, although the mass of Spica is only 2·6 times the mass of the sun.
Finally, in the case of Mizar, we have the sun reduced to a star of
about the seventh, or about five magnitudes fainter than Mizar,
indicating that Mizar is no less than one hundred times brighter than
our sun. These results show the great relative brilliancy of stars with
a Sirian spectrum, when compared with that of the sun, a consideration
which has already been arrived at from other considerations.
CHAPTER V.
VARIABLE AND TEMPORARY STARS.
To ordinary observers, the light of the stars seems to be constant. Even
to those who are familiar with the constellations, the stars appear to
maintain their relative brilliancy unchanged. To a great extent this is,
of course, true; the great majority of the stars remaining of the same
brightness from day to day, and from year to year. There are, however,
numerous exceptions to this rule. Many of the stars, when carefully
watched, are found to fluctuate in their light, being sometimes
brighter, and sometimes fainter. These are known as “variable stars”—one
of the most interesting class of objects in the heavens. Some of these
have been known for a great number of years, and their variations having
been carefully watched, the laws governing their light changes have been
well determined.
We will first consider the variable stars with long periods of
variation, as these generally show the largest fluctuations of light.
Among these, the first star in which variation of light seems to have
been noticed is the extraordinary object, Omicron Ceti, popularly known
as Mira, or the “wonderful” star. It appears to have been first noticed
by David Fabricius in the year 1596. He observed that the star now
called Omicron, in the constellation Cetus, was of the third magnitude
on April 13 of that year, and that in the following year it had
disappeared. Bayer saw it again in 1603, when forming his maps of the
constellations, and assigned to it the Greek letter Omicron, but does
not seem to have noticed the fact that it was the same star which had
been observed by Fabricius seven years previously. No further attention
seems to have been paid to it until 1638 and 1639, when it was observed
at Francker by Professor Phocylides Holwarda to be of the third
magnitude in December, 1638, invisible in the following summer, and
again visible in October, 1639. From 1648 to 1662 it was carefully
observed by Hevelius, and in subsequent years by several observers. Its
variations are now regularly followed from year to year, and it forms
one of the most interesting objects of its kind in the heavens. Its
light varies from about the second magnitude to the ninth, but its
brightness at maximum is variable to a considerable extent. Heis found
its _average_ brightness at maximum in the years 1840–58 to be about the
third magnitude, but on November 6, 1799, Sir William Herschel found it
but little inferior to Aldebaran. On the other hand, at the maximum of
1868, November 7, Heis found it only of the fifth magnitude, and fainter
than he had seen it for twenty-seven years. Sawyer also observed a
maximum of about the fifth magnitude (4·9) on November 10, 1887. M.
Dumenel finds (1896) that in the last twelve periods the magnitude at
maximum varied from 2·5 to 4·7.[115]
It is stated in several books on astronomy, on the authority of
Hevelius, that in the years 1672–76 Mira was invisible at the epoch of
maximum. This is, however, quite a mistake, for it was long since (1837)
pointed out by Bianchi that the supposed non-appearance of Mira in those
years can be simply accounted for by the fact that the star was near the
sun at the time of maxima, and could not be observed. If the star
happens to be at a maximum in April or May, it will be too near the sun
to be seen, and as the mean period is about 331 days, this occurs every
ten years. For this reason the maxima seems to have passed unobserved in
the years 1852, 1853, and 1854, and again in 1883. The star will be very
favourably placed for observation in the year 1897, and some following
years. It has also been stated that Mira wholly disappears at the
maximum, but this is another error, for the star never becomes fainter
than 9½ magnitude at any time, and always remains visible in a 3-inch
telescope. The colour of the star is decidedly reddish, but this hue
seems to be more marked at minimum than at maximum. The spectrum is a
remarkable one of the third type, in which bright lines have been seen
by Espin, Maunder, and Secchi. At the minimum of February, 1896, the
spectrum was photographed by Professor Wilsing, and he found it very
similar to a photograph taken by Professor Pickering some years
previously. The recent photograph shows the lines of hydrogen broad and
bright. There seems to be no other bright lines except those of
hydrogen. The blue end of the spectrum is very similar to that of our
sun, but towards the red end there are “dark flutings, fading towards
the red.” The bright hydrogen lines have only been seen at maximum, but
the instruments used by Professor Wilsing were not sufficiently powerful
to show whether they are also visible at minimum.[116] Professor
Pickering thinks that “probably most of the stars of long period give a
spectrum resembling that of ο Ceti, and having the hydrogen lines G,
_h_, α, β, γ, and δ, bright about the time of maximum. When the
photographic spectrum is faint, only the brighter lines, G and _h_, are
visible.” Within the last few years, Mrs. Fleming, while examining the
photographs of stellar spectra taken for the Henry Draper Memorial, has
detected a number of variable stars of long period by the presence of
bright lines in their spectra. These are mostly telescopic stars.
Although the average period of Mira is about 331 days, it is subject to
marked irregularities, which Argelander has attempted to represent by an
elaborate formula. In recent years, however, the epochs of maxima have
deviated considerably from the dates computed from this formula, and at
the maximum of February, 1896, the star did not reach its maximum light
until nearly two months after the predicted time.
Perhaps the long period variable star next in order of interest—at least
to observers in the Northern Hemisphere—is that known as Chi Cygni. It
was discovered by Kirch in 1686. A mistake is often made about the
identity of this remarkable object It is sometimes confused with the
neighbouring star, 17 Cygni of Flamsteed’s catalogue. At the time of
Flamsteed’s observation, the variable star—which is the true Chi Cygni
of Bayer’s map (made in 1603)—happened to be faint, and Flamsteed, not
being able to find Bayer’s star, affixed the Greek letter χ to his No.
17. It was proposed by Struve to call Flamsteed’s star χ^1, and the
variable χ^2; but there seems to be no necessity to perpetuate
Flamsteed’s error, which has been frequently pointed out. All
authorities on the variable stars now give this variable its proper
designation—χ Cygni. The star varies at maximum from 4 to 6½ magnitude,
and at the minimum it sinks to below the thirteenth magnitude. At some
maxima, therefore, it is easily visible to the naked eye, and at others
it is just below the limit of ordinary vision. At the maximum of 1847,
it was visible to the naked eye for a period of 97 days. The average
period is about 406 days; but, according to Schönfeld—a well-known
authority on the variables—observations indicate a small lengthening of
the period. Observations in recent years show that the minimum occurs
about 185 days before the maximum. This gives 221 days for the fall from
maximum to minimum, and illustrates a feature common to many of the
variable stars, namely, that the increase of light is more rapid than
the decrease. This peculiarity is especially marked in the short period
variables, which will be considered further on. Chi Cygni is said to be
“strikingly variable in colour.” Espin’s observations in different years
show it “sometimes quite red, at others only pale orange-red.” In the
spectroscope, its light shows a splendid spectrum of the third type (or
banded spectrum, very characteristic of these long period variables), in
which bright lines were observed by Espin in May, 1889. One of these
bright lines seems to be identical with the coronal line D_{3}, the
characteristic line of helium.
R Leonis is another remarkable variable star, which is sometimes visible
to the naked eye at maximum. It lies closely south of the star known as
19 Leonis. It was discovered by Koch in 1782. At the maximum, its
brightness varies from 5·2 to 7 magnitude, and at minimum it fades to
about the tenth magnitude. The mean period is about 313 days; but this
is subject to some irregularities, and Chandler finds “good evidence of
cyclical variation of period, with a long term.” The star is red in all
phases of its light, and forms a fine telescopic object. Close to it are
two small stars, which form, with the variable, an isosceles triangle.
The spectrum is a fine one of the third type, a type very characteristic
of these long period variables. Espin finds that the bright bands of the
spectrum are brighter when the star is increasing in light, and fainter
when decreasing. At the maximum of 1889, he found bright lines in its
spectrum.
Another long period variable star which is visible to the naked eye at
maximum is R Hydræ—the Upsilon Hydræ of Bayer—but it is rather too far
south to be well observed in this country. Its variability was discerned
by Maraldi in 1704; but the star was also observed by Hevelius in 1672.
Its light at maximum varies from 3½ to 5½ magnitude, and at minimum it
fades to nearly the tenth magnitude. The period has diminished
considerably since the year 1708, when it was about 500 days. This had
decreased to about 487 days in 1785, to 461 days in 1825, and to 437
days in 1870, and it seems to be still diminishing. Formulæ have been
computed by Gould and Chandler, but do not agree. Schmidt found that the
minimum occurs about 200 days before the maximum. The star is very
reddish, and the spectrum is a fine one of the third type, which Dunér
describes as of “extraordinary beauty,” the typical bands of this type
of spectrum being very large, and perfectly black. At the maximum of
1889, Espin observed a bright line in its spectrum, and finds—as in R
Leonis—that the bright bands are brighter when the star is increasing in
light, and fainter as it decreases.
There is a very remarkable variable star in the Southern Hemisphere
known as Eta Argûs. It lies in the midst of the great nebula in Argo,
and the history of its fluctuations in light is very interesting.
Observed by Halley in 1677 as a star of the fourth magnitude, it was
seen of the second magnitude by Lacaille in 1751. After this, it must
have again faded, for Burchell found it of only the fourth magnitude
from 1811 to 1815. From 1822 to 1826, it was again of the second
magnitude, as observed by Fallows and Brisbane; but on Feb. 1, 1827, it
was estimated of the first magnitude by Burchell. It then faded again,
for on Feb. 29, 1828, Burchell found it of the second magnitude. From
1829 to 1833, Johnson and Taylor rated it of the second magnitude; and
it was still of this magnitude, or a little brighter, when Sir John
Herschel commenced his observations at the Cape of Good Hope in 1834. It
does not seem to have varied much in brightness from that time until
December, 1837, when Herschel was astonished to find its light “nearly
tripled.” He says:[117] “It very decidedly surpassed Procyon, which was
about the same altitude, and was far superior to Aldebaran. It exceeded
α Orionis, and the only star (Sirius and Canopus excepted) which could
at all be compared with it was Rigel, which, as I have already stated,
it somewhat surpassed.”
From this time its light continued to increase. On the 28th December it
was far superior to Rigel, and could only be compared with α Centauri,
which it equalled, having the advantage of altitude, but fell somewhat
short of it as the altitudes approached equality. The maximum of
brightness seems to have been obtained about the 2nd January, 1838, on
which night, both stars being high and the sky clear and pure, it was
judged to be very nearly matched, indeed, with α Centauri, sometimes the
one, sometimes the other, being judged brighter; but, on the whole, a
was considered to have some little superiority. After this, the light
began to fade. Already on the 7th and 15th January, α Centauri was
unhesitatingly placed above, and Rigel as unhesitatingly below, it. On
the 20th, it was “visibly diminished—now much less than α Centauri, and
not _much_ greater than Rigel. The change is palpable.” And on the 22nd,
Arcturus (the nearest star in light and colour to α Centauri which the
heavens afford), when only 10° high, surpassed η, the latter being on
the meridian; η was still, however, superior to β Centauri, α Crucis,
and Spica, and continued so (and even superior to Rigel) during the
whole of February, nor was it until the 14th April, 1838, that it had so
far faded as to bear comparison with Aldebaran, though still somewhat
brighter than that star. In 1843, it again increased in brightness, and
in April of that year it was observed by Maclear to be brighter than
Canopus, and nearly equal to Sirius! It then faded slightly, but seems
to have remained nearly as bright as Canopus until February, 1850, since
which time its brilliancy gradually decreased. It was still of the first
magnitude in 1856, according to Abbott, but was rated a little below the
second magnitude by Powell in 1858. Tebbutt found it of the third
magnitude in 1860; Abbott a little below the fourth in 1861. Ellery
rated it fifth magnitude in 1863, and Tebbutt sixth magnitude in 1867.
In 1874 it was estimated 6·8 magnitude at Cordoba, and only 7·4 in
November, 1878. Tebbutt’s observations from 1877–86 show that it did not
rise above the seventh magnitude in those years, and in March, 1886, it
was rated 7·6 magnitude by Finlay at the Cape of Good Hope. This seems
to have been the minimum of light, for in May, 1888, Tebbutt found that
it “had increased fully half a magnitude” since April, 1887, and might
“be rated as a star of 7·0 magnitude.” From photometric measures made
with the meridian photometer in Peru in the years 1889–91, Professor
Bailey found its mean magnitude to be 6·32, so that probably the star is
now slowly rising to another maximum. Bailey found the hydrogen lines
Hβ, Hγ, and Hδ, bright in the spectrum of its light. Wolf suggested a
period of 46 years, and Loomis, 67; but Schönfeld thought that a regular
period is very improbable. The star is very reddish in colour.
There are many other variables of long period, but they are too numerous
to be described in detail in a work of this character. Particulars
respecting some of them will be found in “The Scenery of the Heavens,”
by the present writer.
We will now consider the variables of short period, which are
particularly interesting objects, owing to the comparative rapidity of
their light changes. The periods vary in length from about 17¼ days down
to a few hours. Perhaps the most interesting of these short period
variables, at least to the amateur observer, is the star Beta Lyræ,
which is easily visible to the naked eye in all phases of its light. It
can be readily identified, as it is the nearest bright star to the south
of the brilliant Vega, and one of two stars of nearly the same
magnitude, the second being Gamma Lyræ. The variability of Beta Lyræ was
discovered by Goodricke in the year 1784. The period is about 12 days,
21 hours, 46 minutes, 58 seconds. At maximum the star is about 3·4
magnitude, and there are two minima, one of magnitude 3·9, and the
other—the chief minima—of 4·5 magnitude. That is, the star has at
maximum 2¾ times the light of the chief minimum, and 1·6 times the light
of the secondary minimum. In other words, if we represent the light of
the star at maximum by 27 candles, placed at a suitable distance from
the eye, the secondary minimum will be represented by 17 candles, and
the chief minimum by 10 candles. These fluctuations, although not very
great, can be easily recognised with the naked eye by comparison with
the neighbouring star Gamma Lyræ. Professor Pickering thought that this
variation in the light of Beta might be explained by supposing that the
star rotated on its axis in the period indicated by the variation, that
the ratio of the axis of the rotating spheroid is as 5 to 3, and that
there is a darker portion at one of the ends, which is “symmetrically
situated as regards the longer axis.” Recent observations with the
spectroscope, however, render this explanation doubtful, and indicate
rather that the star is a very close double or “spectroscopic binary,”
although it does not seem certain that an actual eclipse of one
component by the other takes place, as in the case of Algol. Bright
lines were detected in the star’s spectrum by Secchi so far back as
1866. In 1883, M. Von Gothard noticed that the appearance of these
bright lines varied in appearance, and from an examination of
photographs taken at Harvard Observatory in 1891, Mrs. Fleming found
displacements of bright and dark lines in a double spectrum, the period
of which agreed fairly well with that of the star’s light changes.
Professor Pickering thence concluded that the star consists of two
components, one stellar and the other gaseous, but this conclusion has
been somewhat modified by subsequent investigations. M. Bélopolsky, from
photographs taken with the great 30-inch telescope at the Pulkowa
Observatory, confirms the periodical displacement in the bright spectral
lines “in a period identical with that of the star’s usual double
fluctuation,” but Keeler and Vogel agree that the observed displacements
are incompatible with the supposed occurrence of eclipses. Vogel,
however, is “convinced that Beta Lyræ represents a binary or multiple
system, the fundamental revolutions of which, in 12 days 22 hours, in
some way control the light change, while the spectral variations,
although intimately associated with the star’s phases, are subject,
besides, to complicated disturbances running through a cycle perhaps
measured by years.”[118] The helium line, D_{3}, is visible in the
spectrum.
Another interesting star of short period is Delta Cephei, which is one
of three stars forming an isosceles triangle a little to the west of
Cassiopeia’s Chair, the variable being at the vertex of the triangle,
and the nearest of the three to Cassiopeia. Its variability was also
discovered by Goodricke in 1784. It varies from 3·7 to 4·9 magnitude,
with a period of 5 days, 8 hours, 47 minutes, 40 seconds. The amount of
the variation is, therefore, the same as in the case of Algol, the
star’s light at maximum being about three times its light at minimum.
The period and light curve, however, show, according to Schönfeld, some
irregularities, the computed times of maxima and minima being sometimes
in error to the extent of over an hour. These are, however, small, and,
on the whole, the star seems to be very uniform in its fluctuations.
From seven years’ observations, Argelander found no deviation from
perfect uniformity. The curve representing the light variations is not,
however, very smooth, particularly during the decrease of light, when a
nearly stationary period seems to occur from 16 to 24 hours after the
maximum. The rise from minimum to maximum occupies about one-third of
the period, another example of the feature so characteristic of variable
stars, namely, that the increase of light is quicker than the decrease.
As already stated (Chapter IV.), observations of the spectrum recently
made by M. Bélopolsky, with the great Pulkowa telescope, show that, like
Beta Lyræ, the star is probably a close binary, the period of the
observed fluctuations in the positions of the spectral lines agreeing
with that of the star’s light changes. In this case, however, the lines
are not doubled, as in Beta Lyræ, but merely displaced from their normal
position, indicating that, as in the case of Algol, one of the
components is a dark body. There are, however, no indications that any
eclipse of the bright star by its dark companion takes place. Indeed,
the nature of the light changes, which are continuous and not confined
to a few hours, as in Algol, are inconsistent with the occurrence of an
eclipse. We must, therefore, conclude that the fluctuations of light are
caused in some way by physical disturbances produced by the approach and
recession of the two component bodies in an elliptic orbit round their
centre of gravity. The observations indicate that the component stars,
when furthest apart in their orbital revolution, are separated by a
distance three times as great as when at their point of nearest
approach. The observations also show that Delta Cephei is approaching
the earth at the rate of about 8¾ miles a second. Its spectrum is of the
second or solar type, differing in this respect from the other
spectroscopic binaries, which show a spectrum of the first or Sirian
type. The colour of the star is yellow, and it has a distant bluish
companion of about the fifth magnitude, which may possibly have some
physical connexion with the brighter star, as both stars have a common
proper motion through space.
Another remarkable star of short period is Eta Aquilæ, the variability
of which was discovered by Pigott in 1784. It varies from magnitude 3·5
to 4·7, with a period of 7 days, 4 hours, 14 minutes, but Schönfeld
found marked deviations from a uniform period. It will be seen that the
amount of the light change, 1·2 magnitude, is the same as that of Delta
Cephei. Its colour is yellow, and its spectrum, like that of Delta
Cephei, of the second or solar type. The minimum takes place about three
days before the maximum.
Zeta Geminorum is another variable star with a comparatively short
period. It varies from about 3·7 to 4·5 magnitude, with a period of 10
days, 3 hours, 41½ minutes. Here the variation of light is only 0·8 of a
magnitude, or, in other words, the light at maximum is about double the
light of minimum, as in the case of the Algol type variable, Lambda
Tauri. Its light curve, unlike that of Delta Cephei and Eta Aquilæ, is
nearly symmetrical; that is, the period occupied in the increase of
light is about the same as that of the decrease. Prof. Pickering thinks
that Zeta Geminorum is possibly a “surface of revolution,” one side of
the rotating star being about four-fifths of the brightness of the
other; but Prof. Lockyer finds it to be a “spectroscopic binary,” like
Beta Lyræ and Delta Cephei.
Among variables with very short periods may be mentioned the southern
star R Muscæ, which is close to Alpha Muscæ. It varies from 6·6 to 7·4,
and goes through all its changes in the short period of 21 hours 20
minutes. The minimum takes place about nine hours before the maximum. It
was discovered at the Cordoba Observatory, and Dr. Gould remarks that
“its average brightness is so near the limit of ordinary visibility in a
clear sky at Cordoba, that the small regular fluctuations of light place
it every few hours alternately within or beyond this limit.”
A remarkable variable star of short period was discovered in 1888 by Mr.
Paul in the southern constellation Antlia. It varies from magnitude 6·7
to 7·3, with the wonderfully short period of 7 hours, 46 minutes, 48
seconds, all the light changes being gone through no less than three
times in twenty-four hours! It was for some years believed that the
variation was of the Algol type, but recent measures made at the Harvard
College Observatory show that it belongs to the same class as Delta
Cephei and Eta Aquilæ.
A telescopic variable with a wonderfully short period was discovered by
Chandler in 1894. It lies a little to the west of the star Gamma Pegasi,
and has been designated U Pegasi. It varies from magnitude 8·9 to 9·7,
and was first supposed to be of the Algol type with a period of about
two days, but further observations showed that the period was much
shorter, and only 5 hours, 31 minutes, 9 seconds. The light curve is
quite different from the Algol type, and also from that of Delta Cephei
and other short period variables, the times of increase and decrease of
light being about equal, as in the case of Zeta Geminorum. This fact,
combined with the remarkable rapidity of its light changes, which are
gone through four times in less than twenty-four hours, makes this
remarkable star a most interesting object. Possibly there may be other
stars in the heavens with a similar rapidity of variation which have
hitherto escaped detection.
Several southern variables of short period have been discovered in
recent years by Mr. A. W. Roberts at Lovedale in South Africa.
Unlike the variable stars of long period which seem scattered
indifferently over the surface of the heavens, the great majority of the
short period variables are found in a zone which nearly coincides with
the course of the Milky Way. The most notable exceptions to this rule
are W Virginis with the comparatively long period of 17¼ days, and U
Pegasi, above described, which has the shortest known period of all the
variable stars. Another peculiarity is that most of them are situated in
what may be called the following hemisphere, that is between 12 hours
and 24 hours of right ascension. The most remarkable exception to this
rule is Zeta Geminorum. The above rules do not apply to variables of the
Algol type, which we will now proceed to consider.
Algol, or Beta Persei, is a famous variable star, and the typical star
of the class to which it belongs. Its name, Algol, is derived from a
Persian word, meaning the “demon,” which suggests that the ancient
astronomers may have detected some peculiarity in its behaviour. The
real discovery of its variation was, however, made by Montanari in 1667,
and his observations were confirmed by Maraldi in 1692. Its fluctuations
of light were also noticed by Kirch and Palitzsch, but the true
character of its variations was first determined by the English
astronomer, Goodricke, in 1782. Its fluctuations of light are very
curious and interesting. Shining with a constant, or nearly constant,
brightness for a period of about 59 hours as a star of a little less
than the second magnitude, it suddenly begins to diminish in brightness,
and in about 4½ hours it is reduced to a star of about magnitude 3½. In
other words, its light is reduced to about one-third of its normal
brightness. If we suppose three candles placed side by side at such a
distance that their combined light is merged into one, and equal to the
usual brightness of Algol, then if two of these candles are
extinguished, the remaining candle will represent the light of Algol at
its minimum brilliancy. It is stated in several books on astronomy that
Algol varies to the extent of two magnitudes, but this is quite
incorrect, as a change of two magnitudes would imply that the light at
maximum is over six times the light at minimum, which is more than
double the star’s real variation. The star remains at its minimum, or
faintest, for only about 15 minutes. It then begins to increase, and in
about 5 hours recovers its normal brightness, all the light changes
being gone through in a period of about 10 hours out of nearly 69 hours,
which elapse between successive minima. These curious changes take place
with great regularity, and the exact hour at which a minimum of light
may be expected can be predicted with as much certainty as an eclipse of
the sun.
Goodricke, comparing his own observations with one made by Flamsteed in
the year 1696, found the period from minimum to minimum to be 2 days, 20
hours, 48 minutes, 59½ seconds, and he came to the conclusion that the
diminution in the light of the star is probably due to a partial eclipse
by “a large body revolving round Algol.” This hypothesis was fully
confirmed in the years 1888–89 by Professor Vogel with the spectroscope.
As no close companion to Algol is visible in the largest telescopes, we
must conclude that either the satellite is a dark body, or else so close
to the primary that no telescope could show it. As has been stated in
Chapter III., the motion of a star in the line of sight can be
ascertained by measuring displacements in the positions of the spectral
lines. Now, if the diminution in Algol’s light is due to a dark body
revolving round it, and periodically coming between us and the bright
star, it follows that both components will be in motion, and both will
revolve round the common centre of gravity of the pair. A little before
a minimum of light takes place, the dark companion should therefore be
approaching the eye, and, consequently, the bright companion will be
receding. During the minimum there will be no apparent motion in the
line of sight, as the motion of both bodies will be at right angles to
the visual ray. After the minimum is over, the motion of the two bodies
will be reversed, the bright one approaching the eye, and the dark one
receding. Now, this is exactly what Vogel found. Before the diminution
in the light of Algol begins, the spectroscope showed that the star is
receding from the earth, and after the minimum, that it is approaching
the eye. That the companion is dark and not bright, like the primary, is
evident from the fact that the spectral lines are merely shifted from
their normal position and not doubled, as would be the case were both
components bright, as in the case of some of the “spectroscopic
binaries”—for example, Beta Aurigæ—which has been considered in the
chapter on binary stars (Chapter IV.). Vogel found that before the
minimum of light, Algol is receding from the earth with the velocity of
24½ miles a second, and after the minimum it is approaching at the rate
of 28½ miles a second. The difference between the observed velocities
indicates that the system is approaching the earth with a velocity of
about 2 miles a second. Knowing, then, the orbital velocity, which is
evidently about 26½ miles a second, and assuming the orbit to be
circular, it is easy, with the observed period of revolution, or the
period of light variation, to calculate the diameter of the orbit in
miles, although the star’s distance from the earth remains unknown.
Further, comparing its period of revolution and the dimensions of the
orbit with that of the earth round the sun, it is easy to calculate, by
Kepler’s third law of motion, the mass of the system in terms of the
sun’s mass, and the probable size of the component bodies. Calculating
in this way, Vogel computes that the diameter of Algol is about
1,061,000 miles, and that of the dark companion 830,300 miles, with a
distance between their centres of 3,230,000 miles, and a combined mass
equal to two-thirds of the sun’s mass, the mass of Algol being
four-ninths, and that of the companion two-ninths, of the mass of the
sun. Taking the diameter of the sun as 866,000 miles, and its density as
1·44 (water being unity), I find that the above dimensions give a mean
density for the components of Algol of about one-third that of water, so
that the components are probably gaseous bodies, as Hall has already
concluded.
From the recorded observations of minima in past years, it has been
found that the period of variation of Algol’s light has been slowly
diminishing since Goodricke’s time, and Dr. Chandler finds the present
period is about 2 days, 20 hours, 48 minutes, 51 seconds, or about 8½
seconds less than Goodricke made it. Chandler thinks that this variation
in the length of the period is cyclical, and that it has now about
reached its smallest value, and will soon begin to increase again. He
believes that this variation is probably due to the orbital revolution
of the pair round a third body in a period of about 130 years. M.
Tisserand, however, explains the irregularities by supposing an
elliptical orbit, and a slight flattening or polar compression in the
primary star. Professor Boss is inclined to favour Chandler’s
hypothesis.
It is a curious fact that Al-Sûfi, the Persian astronomer, in his
“Description of the Heavens,” written in the tenth century, speaks
distinctly of Algol as a red star (_étoile, brillant; d’un éclat,
rouge_), while at present it is white, or at the most, of a yellow
colour. A similar change of colour is supposed to have taken place in
the case of Sirius, but the change in Algol seems more certain, as
Al-Sûfi’s descriptions are generally most accurate and reliable.
Stars of the Algol type of variable are very rare objects, only a dozen
or so having been hitherto discovered in the whole heavens. Those
visible to the naked eye, when at their normal brightness, are: Algol,
Lambda Tauri, Delta Libræ, R Canis Majoris, and U Ophiuchi. The
variation of Lambda Tauri was discovered by Baxendell in 1848. It varies
from magnitude 3·4 to 4·2, and its period from minimum to minimum of
light is about 3 days, 22 hours, 52 minutes, 12 seconds. Its
fluctuations have not been so well studied as those of Algol, but it is
known that the “period is subject to marked inequalities,” sometimes
amounting to 3 hours. The variation of light is less than that of Algol,
the light at maximum being only twice the light at minimum. Two candles
at a suitable distance would therefore represent the maximum light, and
one candle the minimum brightness. All the light changes take place in a
period of about 10 hours. The star is white like Algol.
The variability of Delta Libræ was discovered by Schmidt in 1859. It
varies from magnitude 4·9 to 6·1, with a period of 2 days, 7 hours, 51
minutes, 22·8 seconds. The period is, however, according to Schönfeld,
subject to some irregularities. The variation of light is about the same
as that of Algol, the light at maximum being about three times the light
at minimum. The variation takes about 12 hours, of which the decrease
occupies 5½ hours. The star is white like Algol.
The variability of R Canis Majoris was detected by Sawyer in 1887. The
variation is from 5·9 to 6·7 magnitude, or about equal in amount to that
of Lambda Tauri, and the period 1 day, 3 hours, 15 minutes, 55 seconds.
U Ophiuchi was also discovered by Sawyer in 1881. Its variation is from
magnitude 6·0 to 6·7, or slightly less than that of Lambda Tauri, and
the period 20 hours, 7 minutes, 41·6 seconds, but subject to an apparent
diminution. The maximum brightness lasts for about 16 hours, and all the
fluctuations of light take place in the short period of 4 hours. Its
colour is white, like most stars of the Algol type.
U Cephei is a very interesting variable of the Algol type, discovered by
Ceraski in 1880. It varies from 7·1 to 9·5, with a period of 2 days, 11
hours, 49 minutes, 45 seconds. Here the variation of light is greater
than that of Algol, the light at maximum being nearly seven times the
light at minimum. Its rapidity of variation is very great, sometimes
exceeding a magnitude in an hour. The light variations occupy about 6
hours, and the minimum lasts for about an hour and a half, Professor
Pickering thinks that the variation of light is, as in the case of
Algol, caused by an eclipsing satellite, but that in this case the
eclipse may possibly be total, the light at minimum being that due to
the satellite, which may have some inherent light of its own. Lord
Crawford examined the star with the spectroscope, and found that at the
minimum the blue end of the spectroscope faded, and the red was
intensified, which seems to suggest that the light of the star in that
phase shines through a gaseous medium, and that the eclipsing body may
be surrounded with an atmosphere.
Another interesting Algol variable is that known as Y Cygni, which was
discovered by Chandler in 1886, while using it as a comparison star for
the short period variable X Cygni. It varies from 7·1 to 7·9 magnitude,
or about the same amount as Lambda Tauri, with a period of 1 day, 11
hours, 56 minutes, 48 seconds. It has alternate bright and faint minima,
which suggest, according to Dunér, that the star consists of two
_bright_ components, one of them being brighter than the other, and both
revolving round their common centre of gravity in an elliptic orbit,
with a period double that of the light variation. Yendell, who has
carefully observed the star’s fluctuations, fully concurs in Dunér’s
views, and says “the substantial corrections of his fundamental
assumption appears to be proved beyond the possibility of a cavil.”
The variability of the star known as S Cancri was discovered by Hind in
1848. It varies from 8·2 to 9·8, or it is said, at some minima, to 11·7,
with the comparatively long period of 9 days, 11 hours, 37 minutes, 45
seconds. The variations of light occupy about 21½ hours. If the minimum
of 11·7 is correct, we have a variation of no less than 3½ magnitudes,
which implies that the normal light of the star is 25 times its light at
a faint minimum. If this be so, the eclipse must be nearly total.
Argelander found that after the minimum the light increases very
rapidly, and he thinks that the descent from the maximum is even more
rapid.
Some interesting examples of the Algol type of variable have been
discovered in recent years. One detected by Chandler, in 1894, and now
known as Z Herculis, varies from about the seventh to the eighth
magnitude, and has a period of 3 days, 23 hours, 48½ minutes. Faint and
very bright minima alternate in periods of 47 and 49 hours, the ratios
of the light at maximum and minima being 3, 2, and 1. These Professor
Dunér considers, indicate that the star consists of two revolving
components of equal size, one of which is twice as bright as the other,
and he computes that the components revolve round their common centre of
gravity in an elliptic orbit, the plane of which is in the line of
sight, and the semi-axis major about six times the diameter of the
stars. If we assume that the diameter of each component is equal to the
diameter of our sun, I find, from the above data, that the combined mass
of the system is about 1½ times the mass of the sun.
Another remarkable example of the Algol type was discovered by Miss
Wells in 1895. The star lies a little north of the “Dolphin’s rhomb,”
and at its normal brightness is about magnitude 9½. The period of
variation is about four days. The variation somewhat resembles that of U
Cephei. Professor Pickering says: “For nearly two hours before and after
the minimum it is fainter than the twelfth magnitude. It is impossible
at present to say how much fainter it becomes, or whether it disappears
entirely. It increases at first very rapidly, and then more slowly,
attaining its full brightness, magnitude 9·5, about five hours after the
minimum. One hundred and thirty photographs indicate that, during the
four days between the successive minima, it does not vary more than a
few hundredths of a magnitude. The variation may be explained by
assuming that the star revolves round a comparatively dark body, and is
totally eclipsed by it for two or three hours, the light at minimum, if
any, being entirely that of the dark body.”[119] This seems to be an
unique object, and it should be carefully followed through its minimum
with a large telescope.[120]
With reference to the Algol type of variable stars, Chandler finds that
“the shorter the period of the star, the higher the ratio which the time
of oscillation bears to the entire period.” Thus, in U Ophiuchi, with a
period of about 20 hours, the light changes occupy five hours, or
one-fourth of the period, while in S Cancri, which has a period of 227½
hours, the fluctuations of light take up 21½ hours, or only about
one-tenth of the period. In all cases in which the Algol type variables
have been examined with the spectroscope, the spectrum has been found to
be of the first or Sirian type, and they seem to be the only stars with
spectra of the Sirian type whose light is variable. It should be noted,
however, that, on the eclipse theory, the variation of light in these
stars is due merely to an occultation of one star by another, and not to
any physical change in the star itself. The bright star Spica, although
shown by the spectroscope to be a close binary star, like Algol, is not
variable, because, in this case, the plane of its orbit is inclined to
the line of sight, and hence the comparison star does not transit the
disc of its primary. Seen from some other point in space, it would
probably be an Algol variable.
A remarkable peculiarity about the variable stars in general is that
none of them have any considerable proper motion. As a large proper
motion is generally considered to indicate proximity to the earth, we
may conclude, with great probability, that the variable stars, as a
rule, lie at a great distance from our system. In other words, it
appears that the sun does not lie in a region of variable stars, and,
with the exception of Alpha Cassiopeiæ and Alpha Herculis, a measurable
parallax has not yet been found, so far as I know, for any known
variable star.
Plotting the known variables on star charts, I find a marked tendency to
cluster into groups. Thus, in and near the constellation, Corona
Borealis, there are five; near Cassiopeia’s Chair, five. In Cancer there
are four in a limited area. Near Eta Argûs there are several, and in a
comparatively small region in the northern portion of Scorpio there are
no less than fifteen variable stars.
We now come to the interesting and mysterious class of objects known as
“new” or “temporary” stars. These phenomena are of very rare occurrence,
and but few undoubted examples of the class are recorded in the annals
of astronomy. Possibly in some cases they have been merely variable
stars, of irregular period and fitful variability; but others may have
been due to a real catastrophe, such as the collision of two dark bodies
in space, or, possibly, the passage of a bright or dark body through a
gaseous nebula.
The earliest temporary star of which we have any reliable information
seems to be one which is recorded in the Chinese annals of Ma-tuan-lin,
as having appeared in the year 134 B.C. in the constellation Scorpio.
Its position seems to have been somewhere between the stars Beta and Rho
of Scorpio. Pliny informs us that it was the sudden appearance of a new
star which induced the famous astronomer Hipparchus to form his
catalogue of stars, the first ever constructed. As the date of
Hipparchus’ catalogue is 125 B.C., it seems highly probable that the new
star referred to by Pliny was the same as that recorded by the Chinese
astronomer as having appeared nine years previously.
A new star is said to have appeared in the year 76 B.C. between the
stars Alpha and Delta in the Plough, but the accounts are vague.
In 101 A.D., a small “yellowish-blue” star is said to have appeared in
the “sickle” in Leo, but its exact position is not known. In 107 A.D., a
new star is mentioned near Delta, Epsilon and Eta in Canis Major, three
bright stars south-east of Sirius. In 123 A.D., another new star is
recorded by Ma-tuan-lin to have appeared between Alpha Herculis and
Alpha Ophiuchi.
The Chinese annals record that on Dec. 10, 173 A.D., a brilliant star
appeared between Alpha and Beta Centauri in the Southern Hemisphere. It
remained visible for eight months, and is described as resembling “a
large bamboo mat!”—a curious description. There is at present close to
the spot indicated, a known variable star—R Centauri—of which the period
seems to be long and the variation of light irregular. Possibly an
unusually bright maximum of this variable star formed the star of the
Chinese annals, or perhaps the variable star is the remnant of the
outburst which took place in the first century. The variable is a very
reddish star, and at present varies from about the sixth to the tenth
magnitude
A new star is recorded in the year 386 A.D. as having appeared between
Lambda and Phi Sagittarii. Near the position indicated, Flamsteed
observed a star, No. 65 of his catalogue, which is now missing; and it
has been conjectured that the star seen by Flamsteed may possibly have
been a return of the star mentioned in the Chinese annals.
Cuspianus relates that a star as bright as Venus appeared near Altair in
389 A.D., during the reign of the Emperor Honorius, and that he had
himself seen it. There is some doubt, however, about the exact date, as
other accounts give the year 388 or 398. The star seems to have
disappeared in about three weeks.
In the year 393 A.D., another strange star is recorded in the tail of
Scorpio. An extraordinary star is said to have been seen near Alpha
Crateris in 561 A.D. Here again a known variable and red star—R
Crateris—is close to the position indicated by the ancient records.
The Chinese annals record a new star in 829 A.D., somewhere in the
vicinity of the bright star Procyon, and in this locality there are
several known variable stars.
The Bohemian astronomer, Cyprianus Leoviticus, mentions the appearance
of new stars in Cassiopeia in the years 945 A.D. and 1264, and it has
been conjectured that perhaps these were apparitions of Tycho Brahé’s
famous star of 1572 (to be presently described), forming a variable star
with a period of over 300 years. Lynn and Sadler, however, have shown
that the supposed stars of 945 and 1264 were, in all probability,
comets.
Extraordinary stars are recorded near Zeta Sagittarii in 1011 A.D., near
Mu Scorpii in 1203, and near Pi Scorpii on July 1, 1584. It is
remarkable how many of these objects seem to have appeared in this
portion of the heavens.
A very brilliant star is mentioned by Hepidannus as having appeared in
Aries in May, 1012. He describes it as “dazzling the eye.” Other
temporary stars are mentioned in 1054 A.D., near Zeta Tauri, and in
1139, near Kappa Virginis; but the accounts of these are very vague, and
it seems by no means certain that they were really new stars.