FREE ESSAY ON OBSERVING STARS

OBSERVING STARS

Observing Stars
Our view of the sky at night is possible because of the emission and reflection of light.
'Light' is the better-known term for the electromagnetic spectrum, which includes waves
in the visible, ultra-violet, infra-red, microwave, radio, X-ray and gamma-ray regions.
The scale of the spectrum is so large that no region is distinct, several overlap each
other.
Each of these regions in the electromagnetic spectrum represent transverse waves,
travelling as electrical and magnetic fields which interact perpendicularly to each
other, with different ranges of wavelength. The magnetic field oscillates vertically and
the electric field horizontally, and each field induces the other.
By the end of the nineteenth century, Maxwell gave a realistic value for c, the speed of
light: 
c = __1__ = 3 x 108 ms-1 
O(mo eo)
The relationship between the speed of all electromagnetic radiation, wavelength (l) and
frequency (f) is shown to be c = l f.
Because the Universe is so vast, interstellar distances are so great that light emitted
can take upwards of millions of years to reach us. Such large distances are often
measured in 'light-years'; one light-year (ly) is the distance travelled by a wave of
light in a year. Because of the massive speed of light and distances, the light arriving
at us would have left the object many years ago, so that looking at a far away star is
much like looking back in time.
Scientific observation of the stars is difficult because of the distorting effect of the
Earth's atmosphere. One problem is atmospheric refraction-where light is bent. Turbulent
air currents cause varying refractive indices, as there is no uniform air density. This
causes an effect called scintillation, where stars appear to twinkle. The effect on
regions of the electromagnetic spectrum other than the visible part, such as the
absorption of certain frequencies by atmospheric chemicals, and the reflection of waves
by charged molecules in the ionosphere, means that some spectral data is simply invisible
to us on Earth.
The Earth receives electromagnetic radiation of all wavelengths from all directions in
space, but most of the electromagnetic spectrum is blocked out by the atmosphere well
above the Earth's surface, where our eyes and instruments are mostly based. However,
wavelengths from only two regions of the electromagnetic spectrum are able to penetrate
the atmosphere. These two spectral windows in the atmosphere through which we can observe
the Universe are called the optical window-which allows the visible wavelength region
through; and the radio window-which includes the wavelength region from about 1 mm to 30
m. The telescopes used by astronomers on the ground are therefore classed as optical and
radio telescopes. Optical telescopes work by either reflecting or refracting light, using
lenses or curved mirrors to focus the light from a subject to form an image. Radio
telescopes consist of a parabolic reflector and receiver on which the waves are focused.
The gathering and resolving power depend on the diameter of the antenna. Radio
observations are unaffected by the weather or time of day, and because of the larger
wavelength of radio waves, dust in space and atmospheric convection currents are not a
problem. Radio astronomy is used in the chemical analysis of elements (by emission and
absorption spectra); to detect the motion of bodies due to the Doppler effect; and in
investigation into the early Universe and the Big Bang. We can analyse radio waves from
the centres of galaxies, including our own.
Despite the radio window, there are still wavelengths that do not penetrate the
atmosphere. Some radio waves are reflected from the ionosphere, part of the thermosphere,
where streams of charged particles from the sun ionise gas molecules: this is
photo-ionisation. Ultra-violet radiation, X-rays and gamma-rays are also absorbed at this
layer. 
Absorption of the electromagnetic spectrum at various altitudes above Earth occurs to
varying degrees. Much infra-red radiation does not reach ground level because of
absorption in the upper atmosphere by water, and some carbon dioxide and oxygen molecules
that lie between the ground and about 15 km of altitude (the troposphere). Ozone
(tri-oxygen) and di-oxygen in the stratosphere absorbs much of the ultra-violet radiation
(hence the 'ozone layer' at about 30km). A side effect of the ozone layer is that
molecules re-radiate the energy in a few wavelengths of the green, red, and infrared
regions, causing 'airglow'.
It is because of the limitations of Earth's atmosphere, that astronomers learnt the
benefits of observing from beyond it. Placing telescopes and instruments of mountain
tops-to avoid clouds, bad weather and turbulence-or using balloons or aircraft, are
useful, but satellites are far more so. All electromagnetic radiation can be detected,
unaffected by absorption, reflection or refraction, dust, atmospheric haze, airglow,
weather, light pollution or the time of day.
The Hubble Space Telescope is probably the most famous astronomical satellite in orbit
around Earth. Photographs taken by it have far improved detail than an Earth-based
telescope. We have greater knowledge of elements and compounds present thanks to emission
and absorption spectroscopy. The 1983 NASA Infra-Red Astronomical Satellite (IRAS) has
been successful in infra-red observations across the sky, detecting nuclear and chemical
reactions by spectrometry, and hot clusters where stars are born. The 1989 NASA Cosmic
Background Explorer (COBE) satellite undertook a detailed study of background radiation:
the 'echo' of the Big Bang. Low frequency microwaves present today are the result of the
red-shift over a long time of the original, high-energy electromagnetic radiation from
the time of the birth of the Universe. The future of satellite observations lies with
X-ray and gamma-ray astronomy. X-ray images show where high-energy events occur, such as
nuclear processes and matter entering a black hole. Gamma-rays are emitted from only the
hottest and most violent bodies, and although difficult to detect, telescopes are used to
map the Universe.
Most observations surround the light from stars. There are billions of them in the
Universe; we classify stars by their various characteristics. The properties of stars can
be determined by the application of principles explained below. 
All stars visible to us must have surface temperatures high enough to emit light which we
can see from so far away. Some appear brighter than others. The difficulty is in
determining weather a star is very hot and bright, or not as bright but just much closer
to us. We know that very hot things appear 'red hot' or even 'white hot', that the
temperature of an object relates to the colour of light it radiates. The electromagnetic
radiation emitted by any object (whatever its temperature) is known as thermal radiation.
Hot objects such as stars emit high energy, high frequency radiation. At about 1000oc,
thermal radiation falls in the visible region of the electromagnetic spectrum.
To find out the temperature of a star, measurements need to be relative rather than
absolute, as there is no possible way of measuring a star's surface temperature
physically! No object can perfectly emit (or absorb) light in practice, but it is useful
to imagine such a body to make comparisons with: a 'black body'. A black body is a
perfect absorber of light; it follows therefore that it is also a perfect emitter of
light. A perfect absorber would appear totally black; a perfect emitter would emit all
radiation, including visible light, and would appear bright white. We know that a black
body therefore emits a broad range of the electromagnetic spectrum. The most intense
emission will peak at a particular wavelength. The hotter the body, the shorter the peak
wavelength, but the higher the peak. Wein's displacement law states that the peak
wavelength, lmax , is inversely proportional to absolute (actual) temperature of an
object. We assume that a star behaves as a black body. The relationship is shown below:
lmax T = 2.898 x 10-3 m K
Hence, we can relate the colour of a star to estimate its temperature, depending on where
in the electromagnetic spectrum lmax lies. Astronomical objects have peak wavelengths
ranging from radio to X-rays, i.e. surface temperatures from absolute zero to 107 K.
It is apparent that the hotter an object is, the more intense the emission of radiation
is. Luminosity (L) is the total power emitted by a body. The Stefan-Boltzmann law states
that 'the total energy radiated per unit time by a black body is proportional to the
fourth power of its absolute temperature'; it also depends on the surface area (A):
L = s A T4
Stefan's constant (s) = 5.67 x 10-8 W m-2 K-4
The amount of power received per unit area is flux (equal to power / area). Light emitted
from an object spreads out in all directions, the further away it gets the less intense
it becomes according to the inverse square law:
L = d-2
E.g., As Saturn is ten times the distance from the Sun as Earth, the intensity of
radiation is receives is 1/100 th of that for Earth.
The light reaching Earth from the sun can be analysed using a technique called
spectroscopy. It is used to identify the chemical composition of stars (which is mostly
hydrogen and helium), and their surface temperature. Once these are known, stars can be
classified accurately.
An emission spectrum is the spectrum of wavelengths of light emitted from atoms or
molecules. They do this when they lose energy, which corresponds to a specific frequency
of the electromagnetic spectrum. An atom or molecule may become electronically excited,
electrons transfer to higher energy levels, and then later drop back to their normal,
lower energy states, emitting this extra energy as photons of light in the process.
Molecules gain translational, rotational, vibrational or electronic energy, depending on
how much energy they first absorb. They must emit this quantised amount of energy again.
Different elements and have different energy levels, this is why we can associate certain
wavelengths with the physical behaviour of a particular atom. Even small molecules cannot
withstand the high temperatures of stars, their spectra are only visible for cool stars.
An absorption spectrum is apparent when wavelengths of light are missing against the
continuous background of emitted light. These missing wavelengths have firstly been
emitted from atoms in the inner layers of the star, but then absorbed by different
chemicals in the outer layers. Thus we can identify the elements in the outer layers of a
star.
The Balmer series refers to the emission spectrum of hydrogen, specifically for high
energy level electrons dropping back to the second energy level (n=2). Light emitted
falls in the visible region of the electromagnetic spectrum, and the intensity of this
light is an indication of a star's surface temperature. The Balmer series is due to atoms
being excited by kinetic collisions. The electrons of cool atoms occupy their ground
state (n=1), as there are few collisions to excite the electrons. The hotter the atoms,
the more energetic the collisions; more electrons are excited to even higher levels (n=3,
4,.etc). These electrons now absorb wavelengths beyond the Balmer series. The most
intense Balmer emission spectra are from stars with intermediate surface temperatures at
around 10 000K. Most electrons can absorb and re-emit wavelengths of the visible spectrum
at this temperature.
The light from stars travels very great distances, taking a long time, to reach Earth.
Unsurprisingly, it can be affected by the time it reaches us. Of course, our nearest star
is the Sun, and our nearest 'neighbour' is the moon. However, 'near' in space is nowhere
near close enough to actually measure by hand. The first logical estimates used simple
trigonometry in a method called parallax. This is where a distant object will appear at a
different spot when viewed from a different angle. Simply, the position of a star is
measured relative to the background, at the two times when the apparent distance between
these viewing positions is as great as possible. As the Earth rotates around the Sun,
with a radius of 1 astronomical unit (1AU = 1.496 x 1011 m), the greatest possible angle
between two different views of a star is achieved at six month intervals, when the
distance between these two times is 2AU:
The further away the object, the smaller the parallax angle would be, as:
Distance (d) = 1AU
Tan (r)
Distance (d) in parsec (pc) = _____________1_____________
Parallax angle (r) in arc-seconds
Measuring parallax in this way is called annual parallax. It is suitable for objects up
to about a distance of 100pc from us. Earth based instruments are less reliable as the
parallax angle being measured gets smaller, greater measurements have been made by Earth
orbiting telescopes such as 1989 ESA Hipparcus which avoid atmospheric limitations.
We can only estimate the distances of more distant objects such as supernovae. One method
is called spectroscopic parallax, where we can make the assumption that all stars are
equally bright (although we know of course that they are not), and so the brighter a star
the closer it is. 
The apparent magnitude (m) of a star is related to its intensity (I); its is an
observational logarithmic scale. The absolute magnitude is a comparative scale based on
the assumption that all objects are at a distance (d) of 10 pc. The two measurements are
related:
d = 10pc x 10 (m - M) / 5
The distance of an object is related to its intensity (using the inverse square law):
I = L 
4pD2
For objects further away than 10 megaparsecs, astronomers have made use of more unusual
objects as 
reference points in the sky. Cepheid variable stars have luminosity which varies
periodically. They vary in brightness as their surface temperature rises and falls. The
absolute magnitude is directly proportional to the period, and using the above formula
the distance of these stars can be calculated. These stars are present in distant
galaxies, we can deduce how far away these are. Some supernovae behave in the same way.
We know that stars and galaxies are moving away from us, because the spectra lines from
some are shown to have been shifted. This is the Doppler effect, where the spectrum lines
are displaced, because their wavelengths have been changed. The change in wavelength is
related to the velocity:
Df = Dl = v
f l c
The Doppler shift can occur when something is moving towards or away from us, however
receding galaxies is evidence that our Universe is expanding (their light is shifted
towards the red / longer wavelength part of the spectrum). It can also be used to
determine the distance of an object from us. Hubble made the important finding that the
further away a galaxy is, the greater its velocity. Also, all galaxies are generally
moving apart from each other, including ours. Hubble's law depends on the Hubble constant
(Ho), but there is no accurate value for this, due to the inaccurate estimates for
distances by other methods.
v = Ho x d
It is speculated that Ho lies between 40 and 100 km s-1 Mpc-1
The Doppler effect is also used to measure how fast stars and galaxies are rotating, and
the orbital period of binary stars. A pair of binary stars each orbit a common centre of
mass, as they are attracted by each other's gravity. The stars usually have different
masses, and will have different orbits (the radius of which is inversely proportional to
the mass). When the stars are close to each other, it is difficult to distinguish between
them, except by their different spectra (these are spectroscopic binary stars). Each is
identified to be receding or approaching as they rotate, by Doppler shifting.
We can find the combined mass of the two stars (M), based on Kepler's third law of
planetary motion:
M = 4p2r3
GT2
G = Universal gravitational constant = 6.67 x 10-11 N m2 kg-2
The mass of each star can be calculated, as they are known to be in ratio of the distance
to the centre of mass.
We can see that there is so much to be discovered about the sky, over the years
physicists have somewhat overcome the problem of sheer distance across the Universe. We
have catalogued data about many stars, and crucially we can compare other stars to ones
we already know about. We can learn how stars evolve from our observations, however we
can only view a tiny part of history. Star populations are mapped on the
Hertzsprung-Russell diagram, basically a graph of luminosity against surface
temperature:
From it we can examine the life sequences of a star, deduce a star's absolute magnitude,
and then their spectral class according to their surface temperature and other
properties. We can identify what stage in its life a star it. 
Bibliography: 'Astrophysics', Nigel Ingham ; 'Physics Passcards', Graham Booth, Letts