THE COSMIC ENGINE
universe began with a singularity, a point of infinite density.
It is generally believed that our universe resulted from the explosion of
this singularity – the so-called Big Bang.
As a result of the Big Bang, space and energy poured forth from the
singularity and time began.
The universe expanded and cooled and matter condensed out of the mix.
As part of this ongoing process, the Sun and Solar System were formed
over 4 x 109 years ago from a cloud of gas and dust that resulted
from a supernova explosion.
The condensing gas and dust that formed the Sun and the planets contained
all the elements we find today in these celestial bodies.
The planets were formed when matter came together under the influence of
module increases students’ understanding of the history, implications of
Physics for society and the environment and current issues, research and
development in Physics.
Do you live in the Hunter Valley,
Central Coast or Sydney areas??? If so, have you been to Koolang
Observatory yet??? If not, you should go soon!!! Koolang Observatory
is an excellent place to visit for anyone with even a slight interest in
astronomy. Expert guides give talks, show the sun, planets and stars
through large telescopes, explain what you are seeing & answer your
questions. There is also a very good display of many interesting
astronomical phenomena. For more information go to the
Website or phone Koolang on 02 49988216.
GETTING STARTED IN
No matter who you are, or what your background in
Science, your first glimpse of the Moon or Saturn and its rings or Jupiter and
its moons or a Deep Sky object such as the Jewel Box (NGC 4755) through a good
quality telescope is absolutely unforgettable. If you, like many people,
would like to get started in practical astronomy but are not sure how to do so,
then reading this brief section may give you some ideas.
If you are really just interested in some casual
observation and don't want to spend a lot of money, my suggestion is to purchase
a reasonable pair of binoculars. A pair of 7 x 50 or 10 x 50 binoculars
offers a good balance of power, image brightness and light weight. Avoid
fixed-focus or zoom binoculars, which provide inadequate image quality for
viewing sharp pinpoint stars. (Note that a "7 x 50" rating means
7 times magnification and 50 mm aperture for the twin front lenses.) A
good planisphere (such as the Chandler Large Planisphere) and the reference book
"Astronomy Australia" by Dawes et al for the particular year would
also be worthwhile acquisitions.
If you are keen to purchase your first telescope I
would suggest that a Newtonian Reflector on a Dobsonian mount is a good first
choice. If you can afford it a 200mm (8 inch) or 250mm (10 inch) scope is
a very useful size. Both of these scopes are capable of producing very
beautiful, close-up and detailed images of the Moon, planets and many Deep Sky
objects. Both of these scopes are quite easy to re-locate to various sites
as needed. You should be able to get a 200mm Newtonian Reflector on a
Dobsonian mount for around $800. I know it sounds like a lot of money but
it will last you a life time. By way of comparison, it is worth noting
that a similar sized refracting telescope will set you back many thousands of
dollars for negligible improvement in image quality.
If you are serious about purchasing a telescope,
don't go to the local toy store, camera shop or similar! Go to a
specialist - an expert. I can definitely recommend The
Binocular and Telescope Shop, 84 Wentworth Park Road, Glebe NSW 2037, phone 02 9262 1344.
This shop is located on the corner of Wentworth Park Road and Pyrmont Bridge
Road, Glebe. These people are very
experienced astronomers and trained telescope technicians. They know what they are
talking about and are very helpful. The link to their Website is given
below. The address of the Melbourne shop is 519 Burke Road, Camberwell Vic
3124. (Note: I have no financial or business
connection with this shop. It is where I buy all my astronomy tools &
A very useful reference is "The
Night Sky - A Guide to Observing the Sun, Moon and planets" by Steve
Massey (New Holland Publishers, Sydney, 2003).
HISTORICAL DEVELOPMENT OF MODELS OF THE UNIVERSE
For excellent notes on the historical development of models of the universe
from the time of Aristotle to the time of Newton see the following link:
You need to be able to do both of the following:
- Outline the historical development of models of the universe from the time
of Aristotle to the time of Newton. That is, just make yourself some
brief notes that state who did what and when they did it - no details just
the basic facts.
- Choose ONE of the above models. Assess this model to identify
limitations placed on the development of the model by the technology that
existed at the time. For instance, you might choose Kepler's
model. You need to assess its worth - how good was it - what were its
strong points, what were its weak points and then say something about how
the technology of the day or the lack thereof, hindered the development of
the model. For example, think about poor old Kepler. The
development of his model required thousands upon thousands of tedious calculations,
which Kepler did by hand. The lack of
fast accurate calculating machines or computers certainly limited the number
& complexity of the calculations Kepler could make and therefore
ultimately limited the overall extent & accuracy of his model. And
so on. There are other limitations you could give.
Cosmology is the study of the nature and evolution of
the universe. Underlying
current cosmology is the assumption that over very large distances, the universe
is homogeneous and isotropic.
This is called the Cosmological Principle.
To say that the universe is homogeneous means that each region of the
universe has the same properties (eg density, temperature) as each other region.
To say that the universe is isotropic means that the universe looks the
same in all directions.
Up until the 1920’s, it was commonly believed that the
universe was in a steady state of balance.
Galaxies were believed to be stationary relative to one another.
This view was one that dates back to Isaac Newton.
In 1915 Albert Einstein published his Theory of General
Relativity. Today, this is the most
powerful mathematical tool that we have for describing the universe.
It yields field equations which describe the interactions between all of
the matter, radiation and gravitational forces within the universe.
When Einstein first used his equations to describe the
universe he found that they predicted an expanding universe.
As this contradicted Einstein’s expectations, he added a term, called
the cosmological constant (L), to his
equations to counteract the expansion. Other
scientists, notably Alex Friedmann in Russia, Georges Lemaitre in France and
Willem de Sitter in the Netherlands continued to develop the ideas of an
expanding universe using Einstein’s equations without this constant.
This is a good example of what often happens in physics, where theory
works ahead of observation and in so doing helps experimental physicists to
decide the type of observations they should be attempting to make. If observations are found to support the theory, then the
theory becomes more established and accepted by the wider physics community.
In the late 1920’s and early 1930’s, Edwin Hubble and
Milton Humason showed that the wavelengths of the spectral lines from other
galaxies in the universe are red-shifted (ie higher).
This phenomenon is called the cosmological red shift and means that as
time goes by, clusters of galaxies are getting further apart - the universe
is expanding. Hubble &
Humason showed further that the greater the distance to a galaxy the greater the
red shift. So the further away the
galaxy is from us, the faster it is receding.
This universal recessional movement is referred to as the Hubble flow.
The velocities of recession, v, of remote galaxies are related to their
distances from earth, r, via Hubble’s Law:
where H0 is Hubble’s constant (or the Hubble parameter), the
slope of a plot of recessional velocity v’s distance to galaxy.
Up until recently, the actual value of this constant was one of the most controversial issues
in modern astronomy because of its extreme importance – it tells us the rate
at which the universe is expanding. Values
of the Hubble constant ranged from 40 km/s/Mpc to 100 km/s/Mpc depending on the
method of calculation. Results from the WMAP space observatory mission
launched in 2001 to
measure the cosmic microwave background, among
other things, have led to a best estimate so far of the Hubble constant of 70.8 km/s/Mpc,
with an uncertainty of around 5%. (NOTE: 1 Mpc
= 1 megaparsec, a unit of distance = 3.09 x 1019 km.)
Hubble's law suggests that the universe is expanding uniformly.
It is important to realize that each galaxy, or cluster of galaxies, does
not expand as it is held together by its own gravitational forces.
It is the intergalactic space between widely separated galaxies, or
clusters of galaxies, which expands with time.
THE BIG BANG
If we mentally reverse the expansion process, then at some
time in the dim, distant past, all the matter, energy, space in the universe
must have been located in a very small volume – a singularity (point)
of infinite density in fact. It is
generally believed that our universe has resulted from the explosion of this
singularity – the so called Big Bang.
How long ago??? Hubble’s
Law provides a good estimate of the answer.
Let T0 be the time it has taken the most distant galaxies (at
R from us) to move away from us. Then
T0 = R/v = R/ H0R = 1/ H0.
So the reciprocal of the Hubble constant gives us an estimate of the age
of the universe. Try the calculation for yourself. Be careful with
the units. See the
of the Universe Calculation pdf file for the solution.
Measurements from the WMAP
space observatory mission produce a value for the age of the universe of 13.7
billion years with an uncertainty of 1%. More precise measurements
Planck space mission launched by the
European Space Agency (ESA) in 2009 put the age of the universe at 13.8
The Big Bang can be understood as the explosion of space
at the beginning of time. The
four forces (gravitational, electromagnetic, strong and weak) became
distinguishable during the first microsecond.
In the initial stages (first 300 000 years), the universe resembled a
primordial fireball (opaque plasma of electrons, protons and high energy
photons). Plasma is the fourth
state of matter. A plasma is a
mixture of positively charged ions and negatively charged electrons.
It is formed when the temperature is so high that electrons and protons
are too energetic to come together or stay together in the form of atoms.
After about 300 000 years, hydrogen and helium formed and the universe
Note that in standard Big Bang Theory there is no point
asking what conditions were like at or before the Big Bang because time only
began after the Big Bang and matter, energy and space only came into being in
their current, familiar form after the Big Bang. In fact there is a point in time known as the Planck time
(within 10-43 seconds of the Big Bang) when gravity is comparable in
strength to the other forces of nature and due to this the laws of physics as we
know them do not apply. To deal
with the first 10-43 seconds of the Big Bang we need a quantum theory
(As an aside, there are physicists attempting to use String
Theory to extend our knowledge of events back even before the Big Bang.
For instance see “Before the Big Bang” pp.24-27 in New Scientist 3
It is also worth answering a frequently asked question at
this stage. Where did the Big
Bang happen? Can we just trace
the expanding universe backwards and find the place? The answer is no because there is no need!
The reason for this can be explained in many ways but ultimately the
simplest explanation is just this: the Big Bang
happened where you are right now and everywhere else; in the beginning, all
locations we now see as separate were the same location.
Another way to answer the question is to realize that if
the Big Bang brought our entire universe into creation, it could not have
happened at a point inside our own universe, since until the Big Bang happened,
there was no universe. Fascinating stuff, isn't it???
EVIDENCE FOR THE BIG BANG
The early universe was extremely hot – at least as hot as
the surface of the sun. The hot
universe contained many short wavelength photons which formed a radiation field.
As the universe expanded the short wavelength photons had their
wavelengths stretched to become low energy, long wavelength photons.
By the present day, the temperature of this cosmic
radiation field should be quite low, just above absolute zero. This cooled down cosmic radiation field was discovered in
1965. It is called the cosmic
microwave background and has a temperature of about 2.7 K. This microwave background is almost perfectly isotropic,
which confirms Einstein’s assumption that the universe is isotropic.
The Big Bang Theory also predicts the correct proportion of
hydrogen to helium within the universe. Using
the Hubble Law we also obtain an estimate of the age of the universe that is in
good agreement with other age predictions.
MATTER & RADIATION IN THE UNIVERSE
Everything in the universe falls into one of two categories
– matter or energy. The matter
in the universe is contained in such luminous objects as stars, planets and
galaxies, as well as in non-luminous dark matter.
Today when we describe the constituents of matter, we do so
in terms of atoms. The atom
consists of a nucleus containing protons (positive charges) and neutrons
(neutral) surrounded by electrons (negative charges) in orbit around the
nucleus. The nucleus is of the
order of 10-14 m in diameter and the whole atom is about 10-10
m in diameter.
The energy of the universe consists of radiation,
that is, photons. The vast majority
of photons in the universe belong to the cosmic microwave background.
Let us now examine the relative importance of matter and energy in the
Einstein’s mass-energy relationship, E = mc2,
which describes the equivalence of mass and energy, can be used to
express the photon energy in the universe as being equal to a quantity of mass
multiplied by the square of the speed of light. The amount of this equivalent mass in a particular volume,
divided by that volume, is the mass density of radiation, rrad.
Combining E = mc2 with the Stefan-Boltzmann law (see later) we
can calculate the value of rrad
as 4.6 x 10-31 kg/m3.
By determining the total mass of all the stars, galaxies and dark matter
in a very large volume of space and then dividing by that volume, we obtain an
estimate of the average density of matter (rm)
in the universe. This value is
currently thought to be in the range 2 to 11 x 10-27 kg/m3.
This is equivalent to 1 to 6 hydrogen atoms per cubic metre of space, if
the mass of the universe were spread evenly over space.
Since the density of matter in the universe today is much
greater than the mass density of radiation, we say that we live in a matter-dominated
universe. This, however, was
not always the case. As we trace
the evolution of the universe backwards through time towards the Big Bang,
obviously the average mass density of the universe increases, as the volume of
space decreases. At the same time,
the photons become less red shifted and thus have shorter wavelengths and higher
energy than they do today. It turns
out that because of this added energy, the mass density of radiation increases
more quickly than the average density of matter, as we go back in time.
Eventually, we reach a time in the early universe where rrad
Physicists call this state a radiation-dominated universe.
The transition from a radiation-dominated universe to a matter-dominated
one occurred about 2500 years after the Big Bang.
So for the first 2500 years after the Big Bang, radiation
dominated the universe. From that
point in time onwards, as the universe continued to expand and the wavelengths
of the photons became longer and longer, the energy of the photons became lower
and lower and matter gained the upper hand.
As mentioned previously though, this matter was in the form
of a hot, opaque plasma. Then
around 300 000 years after the Big Bang another fundamental change occurred in
the nature of the universe. Around
this time, the energy of the photons became low enough to permit protons and
electrons to combine to form hydrogen atoms.
This is a major step in the evolution of the universe, since today
hydrogen is the most abundant element in the universe.
The epoch when atoms first formed is called the era of recombination.
As hydrogen continued to form, the universe changed from
opaque to transparent, since the hydrogen atoms did not absorb low-energy
photons. This occurred around 1
million years after the Big Bang. Photons
could then travel throughout the universe with little or no interaction with
matter. As a result of this
decrease in interaction between matter and photons, the temperature of matter in
the universe diverged from the temperature of the background radiation.
Today the cosmic microwave background has a temperature of 2.7 K but the
temperature of matter in the universe varies over a huge range from a few tens
of kelvin in interstellar medium to hundreds of millions of kelvin in the
interiors of giant stars.
It is worth noting that although we say that the cosmic
microwave background is isotropic, it is not completely so.
There are tiny variations in the temperature of the radiation field (30 mK)
above and below the 2.7 K average value. These
minute variations indicate that the matter and radiation in the universe were
not totally uniform at the moment of recombination.
These tiny nonuniformities in matter are believed responsible for the
present-day concentrations of mass, such as superclusters of galaxies.
FURTHER DETAILS OF THE EARLY UNIVERSE
By the end of the Planck time, the energy of the particles
in the universe had fallen to 1019 GeV (gigaelectron volts) and the
temperature was around 1032 K. At
that time there was a spontaneous symmetry breaking event in which
gravity was “frozen out” of the primordial fireball.
Gravity then remained separate from the other three fundamental forces.
The universe continued to expand and cool.
At t = 10-35s, particle energy was 1014 GeV, the
temperature was 1027 K and a second spontaneous symmetry breaking
occurred with the strong nuclear force being “frozen out”.
A period of extremely rapid expansion, called the inflationary epoch,
is then believed to have occurred, lasting from t = 10-35s to t = 10-24s.
During the inflationary epoch the universe increased in size by a factor
of about 1050. This exponential expansion is described by
Inflation Theory, an extension to Big Bang Theory that was needed to explain
several aspects of the evolution of the universe that Big Bang Theory could not
At t = 10-12s, the particle energy was around 102
GeV, the temperature about 1015 K, and the electromagnetic force
separated from the weak nuclear force in a third symmetry breaking event.
At t = 10-6s, the particle energy was about 1 GeV and the
temperature about t = 1013 K, low enough to allow quarks to remain
confined to individual protons and neutrons.
This period is called the period of confinement.
(Note that quarks are fundamental particle constituents of protons and
From around t = 10s onwards, deuterium, a nucleus
consisting of one proton and one neutron, began to form but was rapidly broken
back down to its constituents by high energy gamma ray photons. From around t = 3min, the background radiation had cooled
sufficiently to enable the deuterium nuclei to exist freely and to combine with
remaining free protons and neutrons to form helium nuclei.
The process of forming nuclei such as deuterium and helium from protons
and neutrons is called nucleosynthesis.
THE FORMATION OF GALAXIES
It is widely accepted today that there must have been density
fluctuations (slight lumpiness) in the distribution of matter in the early
universe. Otherwise, the
distribution of matter in the universe today would not be as lumpy as it is and
galaxies and clusters of galaxies would not exist.
Through the action of gravity, these fluctuations eventually grew to
become the galaxies and clusters of galaxies that we see today.
In the initial stages of formation, a cloud of gas
consisting of about 76% hydrogen and the remainder helium, probably surrounded
by and interspersed with dark matter (of unknown nature) contains tiny
variations in density. Regions of
higher density will then gravitationally attract nearby material and thus gain
mass. As this contraction occurs,
gas particles collide with each other, the temperature increases and much
hydrogen is ionised. As the
resulting protons and electrons move through the gas cloud they lose kinetic
energy by radiating some energy as photons and by electron collisions with
neutral hydrogen atoms, which result in the atoms absorbing some energy from the
electrons and then re-radiating it as photons.
These two processes allow the gas cloud to cool slightly
and maintain the temperature at around 10 000K.
Since the temperature does not rise, the pressure of the gas does not
increase fast enough to prevent the gravitational forces from collapsing
(crushing) the cloud. The gas cloud
breaks into smaller fragments. The
smaller gas clouds then go through the same process and break into even smaller
Eventually, a stage is reached where the density of the gas
clouds is high enough that photons cannot escape, taking energy with them.
Thus, the gas clouds become opaque to photons, the temperature of the
clouds rise and the pressure forces build to become comparable to the
gravitational forces. At this stage
each small gas cloud enters a phase of very slow contraction and radiates at a
very reduced rate because of the increased opaqueness to photons.
Such an object is called a protostar, and grows in mass as more
matter is accreted from the surroundings.
A protostar will usually spin, due to the residual gas motion.
The matter which accretes to the protostar will form a flattened,
spinning protostellar disc of gas around it.
The slow gravitational contraction of the protostar will
continue until the temperature is high enough for nuclear fusion reactions
to commence at its centre. The
protostar has then become a STAR. The
protostellar disc will continue to undergo its own dynamics, perhaps eventually
producing a planetary system around the star.
The process just described can account for the formation of
a collection of stars – a galaxy -
from a gas cloud. Both
elliptical and spiral galaxies can form in this way.
SPIRAL: If a rotating gas cloud collapses under its
own weight, it will have a tendency to form a flattened disk in a plane
perpendicular to the axis of rotation. If
the rate of star formation is low compared to the time taken for the gas cloud
to collapse, most of the gas will end up in the disc before star formation
occurs. In such a case most of the
stars will form in the disc and a spiral galaxy is formed.
ELLIPTICAL: If the stellar birth rate is high, then
virtually all the gas is used up in the formation of stars before a disk has
time to form. In this way an
elliptical galaxy is formed.
Note that other theories of galaxy formation exist.
The main ones describe the formation of galaxies from the merging of
several gas clouds rather than from the gravitational contraction of huge
isolated gas clouds. Many questions
still exist in galactic evolution. For
instance, why did primordial hydrogen and helium in the early universe clump
into clouds destined to become galaxies? What
is the exact nature of dark matter?
THE FUTURE OF THE UNIVERSE
This section on the future of the universe is not
part of the current syllabus and is included briefly here for completeness.
The future of our universe depends on the average density of matter within
it. The average density of matter,
determines the geometry of the universe and the strength of gravity opposing the
expansion of the universe. Einstein’s General Relativity field equations allow for
three different scenarios:
rm < rc
- this produces
an ever expanding (open) universe
rm = rc
- this produces
a balanced, static state (flat or marginally bounded) universe
rm > rc
- this produces a
universe which will reverse its motion and come to a Big Crunch (closed or bounded
where rc = the critical density of the universe (the
value which when substituted into the field equations yields the flat universe result). Using
a Hubble constant of 75 km/s/Mpc we obtain a
value of 1.1 x 10-26 kg/m3. This density is equivalent to about six hydrogen atoms per
cubic metre of space.
There is growing consensus among cosmologists that the total density of
matter (ordinary matter, dark matter & dark energy) in the universe is equal to the critical density, so that the universe is
spatially flat. This is supported by the data collected by the WMAP space
observatory from 2001 to 2010, which indicates that the universe is flat to
within 0.5% margin of error. The more recent Planck space mission has
produced even more precise measurements in support of a flat universe. It
has provided the most precise measurements to date of the composition of the
universe - 68% dark energy, 27% dark matter and 5% ordinary matter.
Also, in 1998, it was discovered by Brian Schmidt (an Australian), Adam
Reiss and Saul Perlmutter that the expansion of the universe is accelerating.
These Physicists won the 2011 Nobel Prize for Physics for their discovery.
Due to this discovery we now know that our universe will continue to expand at an
ever-increasing rate. See
For further explanation on the future of the universe see
"Foundations of Big Bang Cosmology"
produced by NASA.
OF ASTRONOMICAL DISTANCES:
where d = distance to the star in metres. Note that this is another example of the inverse square law for EM
radiation that we studied earlier. As
we move away from a star (or indeed any light source), the decrease in
brightness is inversely proportional to the square of our distance from the
& ABSOLUTE MAGNITUDES:
Note that in Astronomy we use the magnitude scale to
denote brightness. Hipparchus
invented the scale in the second century BC and modifications over time,
especially in the nineteenth century AD, have produced what we now call the Apparent
Magnitude Scale. Apparent
magnitude is a measure of the light arriving at earth and is directly related to
apparent brightness. Apparent
magnitude describes how bright an object appears to be when seen by an
observer on earth (4).
In the Preliminary Course, we DO NOT need to know the
details of the apparent magnitude scale (or of the absolute magnitude scale). Just
be aware that when reading apparent magnitudes, the greater the apparent
magnitude, the dimmer the star. A
star of apparent magnitude +2 (a second magnitude star) is dimmer than a star of
apparent magnitude +1 (a first magnitude star).
The limit of vision with the naked eye is about apparent magnitude +6.
The Hubble Space Telescope can photograph stars of apparent magnitude +27
(more than 1020 times fainter than the sun) using very long exposure
The absolute magnitude is defined as the apparent
magnitude a star would have if it were located 10 parsecs (32.6 ly) from earth
(4). This quantity measures a
star’s true energy output – its luminosity.
Again the greater the absolute magnitude figure, the less luminous the
star. Absolute magnitude
removes the effect of different distances and allows us to compare the
luminosities of stars with each other. For
instance, the sun appears to be the brightest object in the sky only because it
is so close to us compared to all other stars.
It has an apparent magnitude of –26.8.
If the sun were placed 10 pc from earth, its apparent magnitude would be
+4.8. So, the sun’s absolute
magnitude is +4.8. The most
luminous objects have absolute magnitudes around –10 and the least luminous
around +15. We will do more on
apparent and absolute magnitudes if we do the Astrophysics Option in the
iron bar is heated over an extended period of time, several interesting
observations can be made. While
still at relatively low temperature, the iron bar radiates heat, but no
difference in the colour of the bar is noted. With increasing temperature the amount of radiation that the bar emits
increases very rapidly and visible effects are noted. The iron bar assumes a dull red, and then a bright red colour. As it approaches its melting point the bar becomes a bright
yellow colour. If it could be
raised to even higher temperatures without melting, the bar would become a
blue-white colour. Clearly, with
increasing temperature the bar emits more thermal radiation and the frequency of
the most intense radiation becomes higher.
relationship between the colour of light emitted by a hot body and the
temperature of the body was first recorded by Thomas Wedgewood in 1792. The porcelain-maker noticed that all of his ovens became red hot at the
same temperature, regardless of their shape, size and construction. Experiments by many physicists have since shown that any object with a
temperature above absolute zero (zero kelvin or 0K) emits light of all
wavelengths with varying degrees of efficiency (1).
detailed form of the spectrum of the thermal radiation emitted by most real hot
bodies depends to a certain extent upon the composition of the bodies. Physicists found a particularly useful class of hot bodies to assist with
the study of the relationship between colour and temperature. This type of body is called a
“blackbody” and by definition is
a body whose surfaces absorb all the thermal radiation incident upon them and
allow none to be reflected. All
blackbodies at the same temperature emit thermal radiation with the same
spectrum, independent of their composition. The intensities of the colours in the
spectrum depend only on the temperature (2).
blackbodies are hypothetical, ideal bodies. A perfect blackbody does not reflect any light at all. This is the reason why any radiation that it emits is entirely due to its
temperature (3). Examples of real
bodies which approximate blackbodies include: an object coated with a diffuse
layer of black pigment (such as carbon black); an object containing a cavity
connected to the outside by a very small hole – such an object is called
a cavity radiator (2); stars (including our Sun)
and planets (4). Note that a
blackbody does not necessarily appear black. The Sun does not look black because its temperature is high (around 5800
K) and so it glows brightly. A
room-temperature (around 300 K) blackbody, however, would appear very black (4),
as can be seen from the graph shown below.
The energy density of blackbody radiation inside a
cavity radiator at various temperatures as a function of wavelength is shown
below. Note that the intensity
versus wavelength plot for the radiation emitted from the hole connecting
the cavity to the outside has the same shape. The radiation inside a cavity whose walls are at temperature T has the
same character as the radiation emitted by the surface of a blackbody at
temperature T (2). This
provides a useful means of studying blackbody radiation, since cavity radiators
are convenient to handle both experimentally and theoretically.
mentioned above, the intensity versus wavelength plot for the radiation emitted by a blackbody has the same shape as the energy density versus wavelength
plot shown above. Such a plot
indicates that a blackbody of temperature T emits a continuous spectrum with
some energy at all wavelengths and that this blackbody spectrum peaks at a
wavelength lmax , which becomes shorter with increasing
1893, before plots such as that above had been obtained experimentally, the German physicist Wilhelm Wien
derived a quantitative expression relating the wavelength (or
colour) of the radiation emitted by a hot body to the temperature of the body
(4). To develop this relationship he used an oven with a small hole as a good
approximation for an ideal blackbody. Any radiation that enters the small hole is
scattered and reflected from the inner walls of the oven so often that nearly
all incoming radiation is absorbed and the chance of some of it finding its way
out of the hole again can be made exceedingly small. The radiation coming
out of this hole is then very close to the equilibrium blackbody electromagnetic
radiation corresponding to the oven temperature. Applying his knowledge of
heat and electromagnetism to this situation he came up with what is now called Wien’s
law indicates that as the temperature of a body increases, the wavelength at
which maximum emission of radiation occurs is displaced toward lower
wavelengths. This is in agreement with
the example of the heated iron bar mentioned above, namely that the principal
frequency of the emitted radiation becomes higher (that is, the iron bar changes
colour from dull red to blue-white), with increasing temperature. Clearly,
another way of expressing Wien's law is to say that the wavelength of maximum
emission of a blackbody is inversely proportional to its temperature in kelvin.
Wien’s law is particularly useful for determining the surface
temperatures of stars (4). It is not
necessary to know how far away the star is, how large it is or how much energy
it radiates into space. All we need to know is the dominant wavelength of
the star's electromagnetic radiation. For
example, the star Rigel, 815 light years from earth in the constellation of
Orion, has a continuous spectrum that peaks at a wavelength of 2.87 x 10-7
m in the UV region of the EM spectrum. Its
surface temperature can be found from Wien’s law as:
T = 0.0029 / 2.87 x
TEMPERATURE OF STARS:
from the details above, the surface temperature of a star is related to its
colour. If a star is very hot, its radiation is skewed towards the short
wavelength ultraviolet end of the spectrum. In other words its apparent colour is blue (eg Bellatrix, T = 28 000K). If a star is very cool, its radiation peaks at long wavelengths and its
apparent colour is red (eg Betelgeuse, T = 2400K). Stars can be divided into specific spectral classes that summarize
not only the colour and surface temperature ranges but also the chemical
composition. See the Table below which has been taken from Kaufmann
& Freedman p.470 (4).
atoms especially helium
helium, some hydrogen
hydrogen, some ionized metals
(a Canis Majoris)
ionized metals (Ca, Fe)
neutral & ionized metals, especially ionized Ca
titanium oxide & some neutral Ca
Be aware that astronomers use the term “metals” to refer to any
element above helium in the Periodic Table (4).
Clearly, this is different to the term’s meaning in chemistry.
Note that in researching the colour-temperature relationship for stars you will also
come across something called the B-V colour index (or colour ratio). We will study the details and usefulness of the colour index if we do the Astrophysics Option in the HSC Course. Basically, if the star is hotter than about 10 000 K, it is a very bluish
star with a value of B-V less than 1. If
a star is cooler than about 10 000 K, its B-V value is greater than 1. After measuring a star’s brightness using UBV photometry, an
astronomer can estimate the star’s surface temperature from a blackbody
temperature versus colour index graph.
on the subject of blackbody radiation, it is worth mentioning the Stefan-Boltzmann
Law. This law states that the
energy flux, F (in joules per square metre of surface per second), from a
blackbody is directly proportional to the fourth power of the object’s
temperature, T (measured in Kelvin). Mathematically,
= 5.67 x 10-8 Wm-2K-4, is a constant. Both Wien’s Law and the Stefan-Boltzmann Law are useful tools for
analyzing glowing objects like stars. Note that
you do not need to know the Stefan-Boltzmann Law for the Preliminary Course.
THEIR IMPORTANCE IN ASTRONOMY (NON-EXAMINABLE):
is extension material.
beam of sunlight is shone through a triangular glass prism, the white light is
dispersed, producing a rainbow of colours, which can be displayed on a screen. The rainbow of colours is called a spectrum. In 1814 the German optician Joseph von Fraunhofer discovered that the
spectrum of sunlight contains hundreds of fine dark lines, now called spectral
lines. Fifty years later,
chemists found that they could produce spectral lines in the laboratory and use
these lines to analyze the kinds of atoms of which substances were made.
are three basic types of spectrum:
spectrum – a series of coloured bands
ranging from violet on one end to red on the other. Eg the spectrum given off by a hot glowing body, such as a blackbody
spectrum – a series of bright,
coloured lines on a black background, produced by a hot transparent gas. Eg the spectrum of hydrogen when heated to incandescence by passing an
electric discharge through the gas – a series of four spectral lines (violet,
blue, green & red) is seen on a black background.
spectrum – a series of dark spectral
lines among the colours of the continuous spectrum. Eg the absorption spectrum of sunlight found by Fraunhofer. Note that the dark lines in the absorption spectrum of a particular gas,
occur at exactly the same wavelengths as the bright lines in the emission
spectrum of the same gas.
Spectroscopy is the systematic study of spectra and
spectral lines. Spectral lines are
extremely important in astronomy, because they provide very reliable evidence
about the chemical composition of distant objects, the temperature of objects
and the motion through space of objects.
HERTZSPRUNG-RUSSELL (H-R) DIAGRAM
though in photographs of stars, bright stars appear larger than dim ones, it
would be very wrong to assume that this is actually the case. In fact, to determine the size of a star, an astronomer must
combine information about the star’s luminosity (determined from its distance
& apparent brightness) and its surface temperature (determined from its
spectral type) (4).
the Stefan-Boltzmann law, we can deduce the relationship between a star’s
luminosity, radius and surface temperature:
L = star’s luminosity (watts), R = radius of star (m),
= Stefan-Boltzmann constant and T = star’s surface temperature (kelvin) (4). Clearly then, even a very cool star (low T) can have a very high
luminosity, if its radius is sufficiently large.
relationship between luminosity and surface temperature led the Danish
astronomer, Ejnar Hertzsprung, in 1911 to discover that a regular pattern is
revealed when absolute magnitudes of stars (a measure of their luminosities) are
plotted against their colours (a measure of their surface temperatures). In 1913 the American astronomer,
Henry Norris Russell, independently
found a similar pattern in a plot of absolute magnitudes versus spectral types
(another measure of surface temperature).
Today, the Hertzsprung-Russell (H-R) Diagram is one of the most important in all
astronomy because of its ability to summarize so many trends so succinctly and
because of its usefulness in helping us to understand the evolution of stars. Let us now examine the H-R Diagram in detail.
Hertzsprung-Russell Diagram is a graph of the luminosities of stars against
their colour or surface temperature. On such a graph each data point represents a star whose spectral type and
luminosity have been determined. The
most luminous stars are near the top of the graph, the least luminous stars are
near the bottom. Hot stars (O &
B stars) are towards the left side of the graph while cool stars (M stars) are
toward the right (4). The graph
below shows only a few stars and is designed to show the general patterns which
emerge on an H-R plot.
the different possibilities for axes. H-R
diagrams can be plotted as:
v’s surface temperature (or spectral class or colour index)
magnitude v’s surface temperature (or spectral class or colour index)
of the luminosity relative to the sun v’s surface temperature etc
first important lesson to come from the H-R diagram is the existence of
fundamentally different types of stars. The
plot clearly shows that stars are found in four main groups: Main Sequence,
Red Giants, Super Giants and White Dwarfs.
sequence stars are represented by a band that runs from bright, hot, blue
giant, O class stars in the top left corner down to dim, cool, red dwarf, M
class stars in the lower right corner. Between
80% and 90% of all stars are main sequence stars (1). They are in their hydrogen burning phase and remain on the main
sequence until the hydrogen in their cores is exhausted (5). Note that here the word “burning” is used to mean “nuclear
fusion” – a common piece of jargon used by astronomers. Stars on the main sequence are very stable. The sun is a main sequence star of intermediate luminosity, surface
temperature (and radius) and has sufficient hydrogen to keep it on the main
sequence for at least another five billion years (5).
masses of main sequence stars are directly related to their absolute magnitudes. The brighter the star, the greater the mass of the star; the dimmer the
star, the lower the mass.
Giant stars are found in the upper right
section of the H-R plot. They are
10 to 100 times more massive than the sun and are also about 100 times more
luminous than the sun. Since these
stars clearly have low surface temperature, their high luminosity can only come
from a very large radius. These
stars are cool, reddish in colour and gigantic in size (5). Aldebaran in the constellation of Taurus and Arcturus in
the constellation of Bootes are examples of Red Giant stars.
rare stars are considerably bigger and brighter than typical red giants, with
radii up to 1000 times that of the sun. These
stars are referred to as Super Giants (4). Betelgeuse in Orion and Antares in Scorpius are
two examples of Super Giant stars. Giants
and super giants have thermonuclear (fusion) reactions occurring in their
interiors but the character and location of those reactions can be very
different from those occurring in main sequence stars like the sun (4).
dwarfs form the fourth main group of
stars on the H-R plot and are located in the bottom left section. These stars are about the same size as the earth and can only be seen
with the aid of a telescope (4). They
have temperatures of about 10 000 K on average, are very dim and typically
whitish in colour (5). Note
however, that surface temperatures for white dwarfs do range from 5000 K to 80
000 K and that therefore they can have colours other than white (1). White dwarfs are the remains of giant stars at the end of their life. They have no thermonuclear reactions occurring in their interiors (4).
following section of notes on H-R Diagrams is extension material only and is not
examinable in the current syllabus.
In order to assist with the classification of stars on
the H-R plot, a system of luminosity classes was developed in the
1930’s, based on subtle differences in the spectra of stars. See H-R plot below and Ref.(1) pp.246-250 for more detail.
classes Ia and Ib are composed of Super Giant stars. Luminosity class V includes all main sequence stars. The classes in between provide a useful means of distinguishing giant
stars of various luminosities – Class II Bright Giant stars, Class III Giant
stars and Class IV Subgiant stars. Class
VI consists of Subdwarf stars. There
is no luminosity class assigned to white dwarfs, since they represent a final
stage in stellar evolution in which no thermonuclear reactions are taking place. Consequently, white dwarfs are referred to only by the letter D (for
dwarf). (1 & 4)
describe stars on the H-R diagram by giving both a spectral class and a luminosity
class. The spectral class
indicates the star’s surface temperature and the luminosity class its
luminosity. For instance, Aldebaran
is a K5 III star, which means that it is a red giant with a luminosity around
500 times that of the sun and a surface temperature of about 4000 K. The sun is a G2 V star, which means that it is a main
sequence star of luminosity equal to the sun (obviously) and surface temperature
of about 5800 K. (4)
two-dimensional classification scheme enables astronomers to locate a star’s
position on the H-R diagram based entirely on the appearance of its spectrum. This is very useful, since once the star’s absolute magnitude has been
read from the vertical axis of the H-R diagram, the distance to the star can be
calculated using a method called spectroscopic parallax. We will study this method if we do the Astrophysics Option in the
HSC Course. (See Ref.(1) pp.248-250
for more detail.)
Also see the following for some
good H-R Diagrams:
plot for nearest stars to us
plot for temperature-luminosity information
plot for mass-luminosity information
ENERGY SOURCES IN STARS
The Sun gains its energy from nuclear fusion reactions in which four
hydrogen nuclei (protons) combine to form a single helium nucleus, releasing
energy in the process. The
energy released comes from the conversion of some of the hydrogen nuclei mass
into energy according to Einstein’s equation, E = mc2. The total mass of four hydrogen nuclei is actually greater than the mass
of a single helium nucleus. This
difference in mass is called the mass defect of the helium nucleus and is
the amount of mass converted to energy and given away by the helium nucleus to
stabilize itself during its formation. This
energy given away to stabilize the nucleus is called the binding energy
of the nucleus.
Note that two of the four protons mentioned above actually
decay into neutrons during the fusion reaction. Thus, the helium nucleus
contains two protons and two neutrons.
reactions occur only in the dense, hot core of the sun, since fusion requires
temperatures of around 107 K. Note that for this reason fusion reactions are often referred to as
thermonuclear reactions. Astronomers are
notorious for referring to fusion reactions as “burning”. So they speak of “hydrogen burning” instead of hydrogen fusion and
“helium burning” instead of helium fusion, and so on.
All stars on the Main Sequence
are in their hydrogen-burning phase. Two different fusion mechanisms are responsible for the helium production
and consequent release of energy. For
stars whose core temperatures are above 18 million K, the carbon (or CNO) cycle
is the main mechanism, while for stars whose core temperatures are below this,
the proton-proton chain reaction predominates. We do not need to know the details of these reactions in the Preliminary
When the hydrogen has been exhausted in the core of a Main Sequence star,
hydrogen burning ceases. This
leaves a core consisting almost entirely of helium, surrounded by a shell
through which hydrogen burning works its way outward in the star. The core shrinks and becomes hotter, while the star's outer
layers expand and cool. The result
is a Red Giant star.
When the core temperature
reaches about 100 million K, the fusion of helium (helium burning) begins
there. This process, also
called the triple alpha process, converts helium to carbon and oxygen. In a more massive red giant, helium burning begins gradually whereas in a
less massive red giant, it begins very suddenly, in a process called the
When an old, low mass, star undergoes a helium shell flash, thermal pulses occur
during which more than half the star’s mass may be ejected into space. This exposes the hot carbon-oxygen core of the star. Ultraviolet radiation from the exposed core ionizes and excites the
ejected gases, producing a planetary nebula. (A nebula is an interstellar gas cloud.)
No more nuclear reactions take place in the exposed
core. It becomes a degenerate
(non-contracting), dense sphere about the same size as the earth. It is called a White Dwarf. It
gives off thermal radiation, which causes it to glow. As it cools, it remains the same size but becomes less luminous,
eventually fading into obscurity as a black dwarf.
The fact that energy can be
released from the nuclei of atoms was mentioned above. Energy can be released from nuclei by three processes –
fission (the splitting of heavy nuclei to form lighter nuclei), nuclear
fusion (the joining of light nuclei to form heavier nuclei, as happens in
stars) and natural radioactive decay of the nucleus. Without delving too much into the realms of nuclear physics,
which is another story, it is worth mentioning a few facts about the three types
of natural radioactive nuclear emissions.
Nuclei can emit a,
b or g
These are doubly charged helium nuclei, consisting of a nucleus containing two
protons and two neutrons. Their
double positive charge gives them very good ionizing ability.
Alpha particles therefore interact strongly with matter and have very
poor penetrative power. They are stopped by a few centimetres of air or a few sheets
of paper. Alpha particles are the
most massive of the natural radiations.
These are rapidly moving electrons. Due
to their single negative charge, they have a smaller ionizing ability than alpha
particles and interact less strongly with matter.
They therefore have a higher penetrative ability, being able to pass
through matter hundreds of times thicker than can an alpha particle with the
same energy. b particles can
penetrate about a metre of air.
(g) rays: As
you are already aware, gamma rays are electromagnetic waves.
They have no charge and therefore interact with matter to a much smaller
extent and have a much lower ionizing ability than the other two natural
radiations. Gamma rays have a much
greater penetrative ability than alpha or beta particles and are capable of
penetrating many centimetres of concrete.
It is the relatively larger mass
of the alpha particle that enables it to suffer less deflection as it passes
through both magnetic and electric fields.
VISUAL APPEARANCE OF THE SUN
If you study photographs of the Sun, it appears that the Sun has a sharp well-defined surface. This seems strange, since we know the Sun is gaseous throughout its
volume because of its high internal temperature. The Sun appears to have a surface because all of the visible light
comes from a relatively thin layer of gas called the photosphere. The photosphere is actually the bottom layer of the
and is about 500 km thick (4). The
gases in the photosphere radiate energy with a spectrum approximating that of a
blackbody and with a temperature of 5800 K (4).
Above the photosphere are the chromosphere
and corona, which are transparent to visible light. Below the photosphere is the
solar interior, which consists of the central core with a temperature of about 1.55 x
107 K and two
other regions above the core, which do not concern us here. The total radius of the Sun is 696 000 km (4).
NEVER look directly at the Sun. It
will damage your eyes and can cause permanent blindness. Never attempt to view the Sun with optical equipment (eg binoculars,
telescopes) unless it is fitted with appropriate filters, by someone who really
knows what he/she is doing!!!
FROM THE SUN
The most obvious emissions from
the Sun that reach Earth are light and heat. Other less obvious forms of electromagnetic radiation are emitted by the
Sun and do reach Earth. As
mentioned above, the photosphere radiates like a blackbody with a temperature of
5800 K and therefore emits EM radiation of all wavelengths. The radiation peaks in the
visible region of the
It has been verified by direct
observation from space-craft that charged particles are also emitted from the
Sun and make their way to Earth. A
continuous stream of particles passes the Earth at speeds up to 700 km/s (1). This stream of particles is called the
Solar Wind and consists
mainly of protons and electrons, with some atomic nuclei.
Another type of particle that is
emitted from the Sun and literally passes through the Earth is the neutrino. Neutrinos have no charge, an extremely small mass (almost zero) and
interact only very weakly with matter. Experiments
are in progress today to detect and accurately measure the neutrino flux from
the Sun, in order to gain valuable information about the nature of the Sun’s
core, where the neutrinos originate.
& THE SUNSPOT CYCLE
A sunspot is a dark
region of irregular shape in the photosphere.
Typical sunspots are a few tens of thousands of kilometres across.
They can occur in isolation or in clusters called sunspot groups (4).
Sunspots are not permanent features of the Sun and last from a few hours
to a few months (4).
Using Wien’s Law and
photometric analysis of sunspots, the temperature of the center of a typical
sunspot is found to be 4300 K (4). Clearly,
this is much lower than the average temperature of the photosphere of 5800 K.
In 1908, the American astronomer
George Hale found that sunspots were places in the photosphere where very high
magnetic fields were concentrated. Magnetic
fields of around 0.4 T have been detected in sunspot areas (4), which is about
6000 times more intense than the Earth’s magnetic field measured at its poles.
Thus, sunspots can be thought of
as regions of the photosphere of lower temperature and very strong magnetic
The number of sunspots varies
with a period of about 11 years. This
is called the sunspot cycle. During
the 11-year cycle the number of sunspots increases from a minimum to a maximum
and then decreases again to a minimum. Sunspot
minima occurred recently in 1965, 1976, 1986 and 1996 (4). Sunspot maxima occurred in 1968, 1979 and 1989 (4). This 11-year sunspot cycle is related to a 22-year cycle called the
solar cycle, during which the Sun’s magnetic field reverses.
The sunspot cycle has a direct
impact here on Earth. Sunspot groups are the sources of solar
flares, huge bursts of
energy stretching perhaps 100 000 km in length (1). EM radiation of all wavelengths is released, from
radio waves through to gamma radiation (1). Charged particles are also accelerated to high speeds, many
escaping into interplanetary space as solar cosmic rays (1). These charged particles, mainly protons and helium nuclei,
enhance the solar wind and may reach Earth in 30 minutes or so, disrupting radio communications, causing power blackouts and
difficulties for anyone navigating by magnetic compass and causing auroral
An aurora is a high altitude, natural light display, occurring
most frequently above 60° north or south latitude. It is named according to its location - aurora borealis
(northern lights) or aurora australis (southern lights). Extensive auroral displays are usually accompanied by disturbances in the
Earth’s magnetic field and interference with radio, telephone, and telegraph
Note that there are several links in the Cosmic Engine
section of my Links
Page that provide information useful for answering
Syllabus point 8.5.4 Column 3 dot point 2 on "the effects of sunspot
activity on Earth's power grid and satellite communications".
The remaining sections of notes are extension material only and are not
examinable in the current syllabus.
THE FORMATION OF THE SOLAR SYSTEM
Our Solar System consists
of nine planets: Mercury, Venus, Earth, Mars Jupiter, Saturn, Uranus, Neptune
& Pluto in orbit around the Sun. Many
of these planets have natural satellites (moons). There is an asteroid belt consisting of rocky objects of various
sizes between the orbits of Mars and Jupiter. Beyond the orbit of Neptune there is a belt of asteroids of rock and ice
and short period comets called the Kuiper Belt. Pluto and its moon Charon are actually located on the edge of this belt.
Nine Planets has more information.
Our Solar System formed from
matter created in ancient stars that existed billions of years ago. One or more of these stars ended their life as a supernova (exploding
star) and cast the matter from the star, including many heavy elements,
into the surrounding space. Thus,
the interstellar gas clouds became enriched with heavy elements. One such slowly rotating gas cloud containing hydrogen, helium, traces of
heavy elements, ice and dust particles, contracted under its own gravity, via
the process described previously for the formation of galaxies.
At the center of the cloud, a protostar
formed and the gas cloud flattened into a spinning disc around the protostar. Eventually, after about
100 million years, hydrogen fusion began and the new star (the sun) began to
shine. Over that same 100 million
year period, the solar system formed from the material in the spinning
The four inner planets, the terrestrial
planets, formed through the accretion of dust particles into planetesimals
(small asteroid-like objects with diameters of about 10 km). Under the action of gravity these planetesimals coalesced into protoplanets
(about the size of the moon). These
protoplanets then collided to form the terrestrial planets.
The four large outer planets,
the Jovian planets, began to form by the accretion of planetesimals. Due to the abundance of ice as well as rock grains in the outer solar
system, objects much larger than any of the terrestrial planets formed in the
outer regions. Slow moving gas atoms were captured by the strong
gravitational attraction of these massive rock and ice objects. The accretion of gas happened slowly at first and then more
rapidly as time went by. Eventually,
huge planets formed. These planets
had extremely thick, hydrogen-rich atmospheres, surrounding rocky cores and were
about five to ten times the mass of the earth.
That of course just leaves the
planet Pluto. How did it form? Possibly it began in the same manner as the other Jovian planets with the
accretion of rock and ice to form firstly a planetesimal and then a protoplanet. Perhaps because it was so far out in the protostellar disc, there was
insufficient gas to accrete to this core and so a small, rock and ice mass
formed. At some stage, Pluto was
probably involved in a rather catastrophic collision with another body, which
tore off enough of the planet to form Pluto’s moon, Charon.
EARTH’S MAGNETIC FIELD
Note that the Earth’s internal structure has been shown to consist of:
A solid inner core of iron 1300 km thick
A liquid outer core 2200 km thick, composed of iron with some
A 2900 km thick solid mantle, the upper levels of which are
A solid crust from 0 to 35 km thick
is often convenient to think of the Earth’s magnetic field as being the same
as that of a bar magnet. If the
Earth was in total isolation from the rest of the universe, this may be a
reasonable approximation. Unfortunately,
the Earth’s magnetic field is distorted by interactions with the Sun’s
magnetic field and the solar wind. In
any case, the mechanism by which the magnetic field is produced, cannot be the
same as that which produces magnetism in a bar magnet. The Earth’s interior is too hot for the solid iron core to
be permanently magnetised in that way.
the 1950’s it was suggested that electric currents in the liquid outer core
produced the Earth’s magnetic field.
The molten iron-nickel mixture is stirred into convective motion by heat
generated from residual radioactivity in the core. Stray magnetic fields interact with the moving molten iron to
produce electric currents, which in turn produce their own magnetic fields and
It has recently been shown using super computer
simulations that this theory is probably correct (4).
See any of the
"Magnetosphere" topics at:
for good diagrams of the Earth’s magnetic field.
Earth’s magnetic field carves out a cavity in the solar wind. This cavity is called the magnetosphere.
A shock wave marks the boundary where the supersonic solar wind is very
quickly decelerated to subsonic speeds. The
majority of particles in the solar wind are deflected around the Earth.
existence of the Earth’s magnetic field serves to protect the planet
from the particles of the solar wind, as well as from other ionised cosmic rays.
Instead of striking the Earth’s surface, these particles become trapped
in the dipole field and bounce back and forth along the field lines between the
North and South poles (1). Two
regions of trapped particles (protons & electrons) have been identified and
are known as the Van Allen Radiation Belts.
The inner belt is composed of protons and is at a height of about 4000
km above the Earth’s surface. The
outer belt is composed of electrons at an altitude of about 16 000 km (1).
Particles in the belts that are sufficiently energetic
to enter the Earth’s upper atmosphere near the poles, strike atoms and
molecules there, causing collisional excitation, ionisation and dissociation. When the atoms or molecules recombine, or when electrons drop back down
to lower energy levels, the subsequent emission of light is observed as the auroras,
Earth’s atmosphere consists of approximately 80% nitrogen and 20% oxygen, with
only a small amount of carbon dioxide. The
abundance of oxygen is due to the biological activities of life forms on the
planet. The atmosphere is divided
into four main layers on the basis of temperature differences:
The troposphere – extends from the surface to about 12 km
altitude. Temperature decreases
with increasing altitude within the troposphere. The troposphere contains about 80% of the total mass of the
atmosphere. All of the Earth’s
weather is a result of convection in the troposphere, which in turn is driven by
energy absorbed from sunlight by the Earth’s surface. Features of the troposphere such as water vapour, clouds, air
molecules and dust particles absorb, scatter and reflect back into space,
varying amounts of the in-coming solar radiation.
An important absorption process occurs in the troposphere.
Figures vary, but about 51% of the solar radiation that reaches planet
Earth is absorbed by the surface (land & oceans) and eventually re-radiated
as infrared radiation (heat). Water
vapour and carbon dioxide molecules absorb this infrared radiation. This trapping of energy from sunlight in the atmosphere is
called the greenhouse effect. Without
these heat traps in the troposphere our planet would be about 30o C
cooler than it is at present (4). (Note
that at present there is a very real concern about the fact that CO2
is being added to the atmosphere by combustion processes faster than plants can
remove it. Clearly, the extra CO2
in the atmosphere could have serious consequences for Earth’s future
The stratosphere – extends from 12 km to 50 km above the
Earth’s surface. Ozone
molecules (O3) in the stratosphere are very efficient at absorbing UV
radiation from the Sun. As a
result of this absorption process, heat is transferred to the surrounding air
and the temperature increases with increasing altitude in the
The mesosphere – extends from 50 km to 80 km above the
surface. Very little ozone is
present here and so no UV absorption occurs.
The temperature decreases with increasing altitude.
The thermosphere – extends from 80 km altitude upwards.
Individual oxygen and nitrogen atoms exist here and absorb very
short wavelength UV radiation from the Sun (which oxygen & nitrogen molecules
would not be able to absorb). This
process causes the temperature to increase with increasing altitude.
When an atom in the thermosphere absorbs a UV photon, it can lose an
electron. In fact, the bottom 400
km or so of the thermosphere is the layer of electrically charged particles
(electrons and positive ions) that is responsible for the reflection of radio
waves around the Earth. This
section of the thermosphere is called the ionosphere.
a very good article that presents modern ideas on parallel
universes, see the link below. The article is from the May 2003 edition of
Let me know
via email if this link no longer works. After a certain time articles are
moved to the archive section of the Scientific American site.
Carroll, B.W., & Ostlie, D.A. (1996).
"An Introduction To Modern
Astrophysics", New York, Addison-Wesley Publishing Company Inc.
Eisberg, R. & Resnick, R. (1974).
"Quantum Physics of Atoms,
Molecules, Solids, Nuclei and Particles", Canada, Wiley
Playoust, D.F., & Shanny, G.R. (1991).
"An Introduction to Stellar
Astronomy", Queensland, The Jacaranda Press
Kaufmann, W.J. III, & Freedman, R.A. (1999).
Edition), New York, W.H. Freeman & Company
Bhathal R., (1993). "Astronomy
for the Higher School Certificate", Kenthurst, Kangaroo Press Pty Ltd