An Introduction to the Universe

Exercise:

You have 10 minutes and a single sheet of A4 paper to draw a scale model of the solar system using the following values:

How did you do? http://www.exploratorium.edu/ronh/solar_system/

Task

Do a research project to investigate and present aspects of the history of astrophysics. You should pick one of these titles and write a presentation to other students of your level lasting no longer than 5 minutes.

• Ptolemy and the geocentric model

• Copernicus

• Galileo, Kepler and Newton

The essential points that should have been mentioned.

Task:

Complete the Introductory Astrophysics questions, to compile a set of notes on the the basic nature of the Universe.

Bishop pp 1-8, Kirk and Hodgson pp294-295, Kirk pp120-121

Back to Astrophysics Option

Astronomical Distances

The Distances of Stars

Of the objects that we see in they sky which we call 'stars', some are stars and some are galaxies or clusters of galaxies. The stars tend to be within our own galaxy -the Milky Way, and the other galaxies and clusters are extra galactic.

Galileo's observations of small, feint, stars with his telescope, and Newton 's idea of an infinite universe tended to point to the stars being different distances away. Neither was able to measure the distance to the stars though.

Their best effort was to classify them according to their ' apparent brightness'. If, they assumed, all stars are the same brightness, then the apparent differences that we see are due to their differing distances from us. Therefore the brighter the star the nearer to us it is.

The inverse square law was used to predict relative distances. ie/ A star 4 times as bright as another must be 2 twice as close to us.

Question:

If Newton estimated the Sun to be a million million (10¹²) times brighter than Sirius (the brightest star in the night sky), how far away did he reckon that Sirius was?

Measurements of Astronomical Distances

We measure distances in terms of:

The Astronomical Unit (AU)- the distance from the Earth to the Sun (1.5 x 10¹¹m)

The light year (ly)- the distance light can travel in one year (365x24x60² x 3x10⁸= 9.5 x10¹⁵m)

The parsec (pc)- explained below...

Parallax

Consider the play excerpt about Galileo.

This lends itself to the idea of parallax as an observational tool to indicate movement within an external frame of reference, ie/ relative movement.

We measure parallax using degrees of arc.

• The sky overhead forms a dome of 180^o.

• Divide the sky into 180 equal sections. Each one represents one degree of arc (finger nail ~ 1^o, hand ~ 10^o)

• Then divide each degree into 60 minutes, and again into 60 seconds.

• We usually use arc seconds to describe the observations of parallax that we make.

Using Parallax to measure the distance to a star

Parallax is half the angle through which a star's direction changes as the Earth moves from one extremity of it's orbit to the other.

It is important to have a constant against which to calibrate the angle of a star when measuring its parallax, and to do this we use a very distant star.

The distance of a star in parsecs is.

Distance (pc) = 1/ parallax (s)

If the parallax of 61 Cygni is 0.3 seconds of arc, then its distance is.

= 1/ 0.3 = 3.33 pc from Earth.

Question:

What constraints can you see with using the stellar parallax method for measuring the distances to stars? What is the limit of using this method?

Tasks:

• Find out the sizes of the following distances in AU, ly and pc, compare and contrast them in an appropriate table: Radius of the solar system, Average distance to our 5 nearest stars (what are they?), radius of the Milky Way galaxy, distance to our 5 nearest neighbouring galaxies (what are they?), estimated radius of the universe.
• Use the parallax measurement system to estimate the angular distance between two stars visible in the sky at the moment. Also find out about the nature of each of the stellar objects mentioned on the map: quazars, cepheids, open and globular clusters, nebulae, etc.
• Attempt planning 1 from the practical exercises

Bishop pp50-53, Kirk and Hodgson pp296-301, Kirk pp122, 126-128

Back to Astrophysics Option

The Magnitude of Stars

The Cepheid Method

Using the paralex method for measuring the distance to stars works well for stars that are relatively close to us, but the method breaks down for more distant stars (several hunderd parsecs).

Question:

What is the furthest star that we can measure from Earth, if the maximum parallax resolution we can measure is 0.01"?

For stars more distant than this we use the Cepheid method.

A cepheid variable, is a star whose intensity varies in a regular and constant way.

Observations show that the maximum intensity of these stars is directly related to their period of variation.

Luminosity and Brightness

We measure the brightness of stars in several ways...

Luminosity is the power emitted by a star (how bright the star really is), brightness is the power recieved per unit area to an observer on Earth (how bright it appears to us).

Remembering the inverse square law, the two quantities are related by the equation...

The Magnitude Scale

How bright a star appears is dependant on two factors:

• How bright the star is
• How far away the star is

The apparent magnitude(m) of a star tells us about how bright that star appears from Earth.

But in order to gain a real idea about a star we need to standardise the scale to account for the different distances of different stars.

The absolute magnitude(M) scale accounts for this and supposes all stars to be a distance of 10 parsecs away. Therefore the absolute magnitude of a star is how bright it would apear if it was viewed from a distance of 10pc.

Classically the brightest stars visible to the naked eye had an apparent magnitude of 1 and the dimmest an apparent magnitude of 6.

When measurements could be made it was found that there was a difference in intensity of 100 times between a magnitude 1 and 6 star.

2.5 x 2.5 x 2.5 x 2.5 x 2.5 = 2.5⁵ = 100

Therefore a jump of 1 magnitude represents an intensity increase of 2.5 times.

Doing Magnitude Calculations

We calculate apparent magnitude, m , of a star from its measured brightness, b, using...

m = -2.5 log b + constant

Now comparing two stars of different apparent magnitudes, A and B.

m_A - m_B= (-2.5 log b _A+ const) - (-2.5 log b _B + const)

= -2.5 (log b _A- log b _B)

= -2.5 log b _A / b _B

Assuming we are comparing the apparent magnitude, m , with the absolute magnitude, M , of a single star.

M - m = -2.5 log (b₁₀ /b)

Where b ₁₀ is the brightness of the star at a distance of 10 pc, and b is the observed brightness of the star.

But the inverse square law gives us b₁₀ /b = (d/10)² (since b is proportional to 1/d² from the inverse square law).

So...

M - m = -2.5 log (d/10)² = -5 log (d/10)

M = m - 5 log (d/10)

Questions:

•  A certain type of variable star are known to have a average absolute magnitude of 0.0. Such stars are observed in a particular star cluster to have an apparent magnitude of +16.0. What is the distance to the star cluster?

•  The star Procyon in Canis Minor (the Small Dog) is a prominent star in the winter sky, because it's apparent magnitude is +0.37. It is also one of the closest stars, being only 3.51 parsecs from Earth. What is the absolute magnitude of Procyon? How many times brighter or dimmer than the Sun is it?

•  Barnard's star can only be seen with a telescope because its apparent magnitude is +9.54. Its distance from Earth is only 1.81 pc. How much closer to the Earth would Barnard's star need to be for it to be visible to the naked eye?

Back to the Cepheid Method

We can measure the distance to cepheid variable stars by the finding their period and hence absolute magnitude, and comparing this and their aparent magnitude to deduce their distance.

Other stars distances can be established if they are close to cepheid variable stars. Similarly the distances to galaxies can be determined if they have cepheid variable stars in them.

Bishop pp53-62, questions 1-5 pp83-84

Kirk and Hodgson pp296-301, Kirk pp122, 126-128

Back to Astrophysics Option

Fusion and Stellar Evolution

The Sun and indeed all stars release energy through a process of nuclear fusion* occurring in their cores.

*Fusion is the opposite process to that of fission, where the nucleus of a large unstable atom is split releasing energy. Note when the opposite process occurs with atoms smaller than Iron energy is also released.

The fusion process in a young star (made mainly of Hydrogen) is.

In the fusion process, when the 4 H nuclei are converted into a He nucleus a small amount of mass is lost as energy.

Question:

Complete the worksheet to find out how much energy is released by the Sun.

Evolution of Stars

We have so far considered the 'main sequence lifetime' of a star. This is the period of its life while it undergoes hydrogen burning. During this period the star will reside on the main sequence of the HR diagram- see next lesson.

Once a star has burnt all (or most) of its hydrogen, then depending on several factors (most importantly it's mass) it will evolve into different forms.

Typically a star will expand to become a red giant or supergiant star. This happens because the star starts to fuse heavier elements at its core (Helium, Silicon, etc). The requirements for this are higher temperatures and pressures, leading to greater outward radiation pressure and consequent expansion.

Task:

Find details about the further nuclear fusion processes, specifically the CNO cycle (Bishop pp143).

What happens next?

Eventually the star will run out of fuel that it can burn to power itself. This is a critical point and what happens next is mass dependent.

The Chandraskahar Limit is a mass limit placed on stars which will determine their evolutionary path. The value of the Chandraskahar Limit is about 1.4 solar masses.

Below the Chandraskahar Limit

A star of mass less than 1.4 solar masses will contract once the fusion processes at it's core have finished, since the gravitational force is no longer balanced by radiation pressure. It will contract until its collapse is prevented by electron degeneracy pressure. At this point it is a small dense hot object and then continue to cool. Initially, while still hot it will be a white dwarf, but as it cools its colour will change.

Above the Chandraskahar Limit (1.4M_o<M<8M_o)

In this range the same process as above happens, but rather than stopping at a white dwarf, the gravitational force continues to contact the stellar remenant until the pressure within it is so large that it forces the electrons and protons in individual atoms togther to form neutrons (inverse beta decay). Neutron degeneracy pressure then prevents any further gravitational collapse. Still though we have a very hot, very dense, very small remenant called a neutron star.

Above the Chandraskahar Limit (M>8M_o)

For truly massive stars a very rare and special process occurs. A star is thought to go 'supernova' at masses above 8 solar masses. At this point the energy of the neutrinos released in the inverse beta decay is go great because of the shear rate of contraction, that the collapse causes an explosion of huge magnitude (10⁴⁴J). This is how all the elements heavier that Iron must originally have been made!

The Tolman-Oppenheimer-Volkoff theory of stellar structure

What is left after the supernova- the supernova remnent (SNR) could then take one of 3 forms: white dwarf, neutron star or black hole, depending upon its mass.Tolman-Oppenheimer-Volkoff theory of stellar structure places mass limits on the SNR, and defines which form it will take.

Task:

Find out about and write a page on the nature of pulsars and quasars.

Bishop pp158-182, Kirk and Hodgson pp302-303, Kirk pp132-133

Back to Astrophysics Option

The Hertzprung-Russell Diagram

At a very basic level the HR Diagram allows us to classify stars. At this level we use the HR diagram as a graph of Absolute Magnitude plotted against Temperature.

Task:

Plot a HR diagrams for the following data: HR diagram plotting exercise

Conclusions:

• The sun as a very average star is right in the middle of a HR Diagram:

• The main sequence is a line running diagonally across the diagram and on it most stars lie during the Hydrogen burning phase of their lifetimes.

The other main types of stars that we need to add to the HR Diagram are White Dwarfs, Red Giants, Hot Blue stars and Cool Dim stars.

The two main questions that will be raised from this is .

1. Why do we call some stars red, some blue and others white?
2. How do we know what temperature a star is to be able to place it in the right place on the HR Diagram.

Different colours occur because of different wavelengths of light. This is the principle behind knowing the colour of a star.

We can measure the wavelength of light emitted from a star and from it deduce the colour in the e-m spectrum that this corresponds to. In this way we classify stars into 'Spectral Types'.

Mnemonics to remember this by

Therefore the temperature of a star can be inferred from the wavelength of the radiation it gives out.

All stars are black bodies, meaning they emit e-m radiation at all wavelengths - they are (near to) perfect emitters.

The radiation profile from different stars is different - they give out different amounts of radiation at different wavelengths depending on their temperature. The peak wavelength is the wavelength at which the highest intensity of radiation is emitted. This is used to deduce temperature using 'Wein's displacement law'.

Wein's displacement law:

l_max T = constant

The constant has the value 2.9 x 10^-3mK. This assumes the wavelength to be in metres and the temperature to be in Kelvin.

The Stefan-Boltzmann Law

We can go further now and calculate the radius of a star using the Stefan-Boltzmann Law:

L=AsT⁴

Knowing (or having calculated) the luminosity and temperature of a star we can use the above relationship to determine the surface area of the star, and thus its radius.

Now have a go at the HR diagram worksheet

The HR diagram and stellar evolution

Task:

Draw the evolutionary paths for the following stars on an HR diagram.

• A low mass star (M<1.4M_o)
• A high mass star (1.4M_o<M<8M_o)
• A high mass star (M>8M_o) whose SNR mass exceeds the Tolman-Oppenheimer-Volkoff limit.

IMPORTANT NB: A common misconception is that stars move along the main sequence during their lifetimes. THEY DON'T (much)!

Bishop pp146-157, Kirk and Hodgson pp297-298, 302-303, Kirk pp125, 132-133

Back to Astrophysics Option

Stellar Spectra

Revision:

Atoms are made up of a small dense nucleus and a relatively sparse region in which the electrons orbit. The electrons in an atom absorb and emit energy.

Initially an electron may be in a ground (unexcited) state. Given a certain amount of energy though and the electron will become excited and 'jump' to an excited state. The electron can then absorb more energy and jump 'up' to a higher excited state. To jump 'up' a level requires the electron to absorb a certain 'discrete' amount of energy, because these levels only exist at certain energies.

This explains why certain atoms only absorb radiation of a particular wavelength, because a quantum of radiation at this wavelength is enough to raise an electron up an energy level.

In a solid and a liquid the forces between atoms are complex, and the energy levels can get mixed up. However in a gas the atoms are well spaced and this doesn't happen. Therefore we are able to recognise a gas from the specific wavelengths of radiation it absorbs. This is called its absorption spectrum.

We can identify gases from the missing radiation in their spectra!

When the electrons in an atom become 'de-excited' they emit the energy they absorbed to get excited and fall back an energy level. Since the energy levels are discrete the photons the electrons emit are the same wavelength as those that they absorb. The spectrum of photons emitted from a gas is called its emission spectrum.

So if a gas absorbs radiation of a particular wavelength, but then emits it again, how do we know that this is happening?

The atoms in a gas absorb the radiation as it travels in a particular direction. But when the atom emits the radiation, it does so in all directions. Thus less radiation of the absorbed (and emitted) wavelength ends up travelling in the original direction and hence less ends up at the observer.

So...

Dark lines in the absorption spectrum of any element exactly match the bright lines in its emission spectrum.

Eg, Helium.

Looking at the absorption spectrum of sunlight, and comparing with the emission spectrum of Iron in the same range, proves its presence.

Task:

Explain why a sodium vapour lamp, a mercury vapour lamp and a neon sign each produces its own distinctive colour of light.

End Note: Radiation reaches Earth, not just as light, but as all forms of electromagnetic wave. We are able to learn just as much from the other types of radiation and by interpreting their data as we are from light.

Spectroscopic Parallax to Find the Distance to Stars

We can gain a lot of information about a star by looking at its spectral lines.

We can determine its composition from the relative intensities of its Hydrogen and Helium lines (and higher elements too).

Stars are generally black body emitters (emit radiation at all wavelengths) but by Wein's Law the temperature of the star determines the peak wavelength emitted. Studying this peak tells us about the absolute magnitude or luminosity of the star (from the HR diagram).

To determine the distance to the star is then a simple matter of comparing the inferred luminosity (or absolute magnitude) with its brightness (or apparent magnitude).

Theoretically there is no limit to the distances we can measure using this method, however practically we a limited to distances of about 10 Mpc. Beyond this distance our inferrences of absolute magnitude are very inaccurate.

Questions on Astronomy

Answer the past exam questions on the course so far.

Bishop pp86-98, questions pp99-100, Kirk pp123

Back to Astrophysics Option

Galactic Distribution and the Cosmological Principle (Higher level only)

Galaxies

We know that stars are not just scattered randomly around d the universe, but are grouped into clusters known as Galaxies. This theory was developed first by the observation and deduction of the presence of the Milky Way (Galaxias is Greek for milk).

The Milky Way

You can see a faint broad band of light running across the sky if you look up on a clear moonless night somewhere where there is little 'extinction' from street lights. Galileo observed that this was in fact a collection of faint stars which make up our local galaxy.

Star counts conducted by Herschell suggested that there were similar numbers of stars in all directions. This led to the theory that the Sun is situated close to the centre of the Milky Way.

But...

Early this century the distances to Cephid Variables that surround the galaxy were measured and the results suggested that the Sun was in fact much closer to the edge of the galaxy than was previously thought. Incidently the same method gave the diameter of the galaxy as 100 000 light years.

Diagram of globular clusters surrounding the Milky Way

The two theories were reconciled when it was realised that large amounts of dust at the centre of the galaxy were obscuring the light from more distant stars.

Studies, of these areas of gas which occur throughout the galaxy, in the 1950s showed that the Milky Way is a spiral structure. This was deduced by observing the Doppler shift of the 21 cm radio waves emitted from the dust clouds.

Other Galaxies

Immanuel Kant suggested that bright patches of light (nebulae) may be other galaxies in 1755. Hubble confirmed this idea in 1924 when he discovered Cephid Variables in other nebulae and spiral structures of stars using a 100 inch telescope. These Cephids turned out to be much further away than stars in the Milky Way.

Hubble calculated the distances to these galaxies and was able to map them in relation to the Milky Way. He was also able to work out their absolute magnitudes and discovered that they varied by only about 2.5.

Despite not being able to pick out Cephid Variables in very distant galaxies, by assuming an average brightness Hubble was able to determine their distance from their apparent magnitudes. Thus Hubble mapped an area of the Universe to a distance of 500 million light years.

The map showed that on a large scale the universe is isotropic (the same in all directions). This is called the cosmological principle .

Hubble effectively completed the Copernican Revolution. The Earth far from being at the centre of the Universe (as per the Ptolemaic model) it is but a satellite of a small star in an average galaxy far from the centre of the millions of known galaxies.

The difference between the apparent magnitude and the absolute magnitude of a Cephid variable in the Andromeda galaxy is 24.

•  Calculate the distance to the galaxy.

•  Compare the distance to Andromeda to the diameter of the Milky Way.

Task:

Extension: P/copy of p24 read Box 2D More about galaxies, and answer SAQ2.9 and 2.10.

Video:

Watch the TED talk given by Prof George Smoot, University of Berkely, Calif, on the Design of the Universe.

Bishop pp15-18, Kirk pp134

Back to Astrophysics Option

Olbers' Paradox and Hubble's Law

Newton's Model of the Universe

What conditions did Newton's theory of gravitation put on the Universe, given that Newton saw the position of stars in the night sky remain fixed?

• uniform
• infinite

Olbers' Paradox

Newton 's model of an infinite universe was to an extent confirmed by Hubble's observations. Stars were grouped in to galaxies, and clusters of galaxies. But these clusters seemed to be spread out uniformly in space as far had been mapped (500 million light years). There seemed no reason to assume that the universe didn't continue in this uniform static way forever. Thus a Static, infinite uniform universe.

BUT: Why isn't the night sky as uniformly bright as the surface of the Sun?  If the Universe has infinitely many stars, then it should be.

This can now be rejected on many levels: See handout: Olbers Paradox .

The main two assumptions of Olbers' that reject this paradox are:

• That space extends infinitely in all directions.
• That the universe is static.

Why must a static universe also be infinite?

Hubble's Law

Hubble noticed that the spectra from the galaxies he was observing were red-shifted by the Doppler effect.

Hubble noted this and determined that the speed of the galaxies he was observing was.

and because the galaxies were red shifted he could further say that these galaxies were receding at this speed.

Questions (1st from handout)

Hubble's Law and the Big Bang

Hubble measured the recessional speed of many galaxies by measuring their red-shifts.

He plotted recessional speed against distance to the galaxy and found a direct proportionality.

This gives us Hubble's Law:

The recessional speed of a galaxy is directly proportional to its distance.

and this gives us solutions to Olbers' paradox...

• Firstly that if galaxies are receeding fron us then we receive less rediation from them each second. So if a galaxy is receeding from us and is a long way away then we receive very little radiation from it.

• Secondly that if galaxies are moving quickly away from us then the radiation that we receive from these galaxies will be highly red-shifted and the energy of each photon reaching us will be smaller (E = hf and f is small, since l =c/f and l is getting larger with red-shift). Still however there should be an omidirectional source of low energy photons reaching us from these receeding galaxies.

Doesn't this mean we are at the middle of the universe though? No!

Hubble's observations led to the big bang theory. If all the observable galaxies are at currently moving away from us, then if you play time backwards then they would all be moving together.

Demo with balloon

This would mean that at some point in the distant past all the matter in the universe existed at the same place, and an explosion must have occurred in order to cause it to disperse. Thus the big bang .

In addition to this Hubble had also worked out a way to determine how long ago the big bang happened.

Given that speed = distance / time,

then time = distance / speed.

For a typical galaxy 5x10⁹ light year distant, and receding at a speed 0.25 times the speed of light, then the time since the Earth and it were in the same place is.

= 5x10⁹/ 0.25

= 2 x 10¹² years (or 20 billion years)

This is the value for one galaxy only though and other galaxies reveal slightly different results for various reasons. The most accurate value we can get is an average from all the galaxies we can observe, and this turns out to be 1/ gradient of the Recessional speed - distance graph that Hubble originally plotted.

The gradient of this graph is called the Hubble constant, and its exact value is one of the most hotly debated topics in astrophysics today. It is generally agreed to be somewhere between 50 and 75 x 10^-12y^-1 . Which puts the age of the universe at about 12 2 billion years.

A few more questions...

1) Why can the age of the universe be calculated from any point on the Hubble graph?

2) How can the expanding universe provide solutions to Oblbers' paradox?

3) Astronomers quote the Hubble constant as 80km s^-1Mpc^-1, so what is its age?

Task:

Plot a Hubble Graph with the data in the worksheet for different galaxies to determine the age of the Uninverse.

Find out about how the galaxies in the observable universe are thought to be distributed. What is meant by the terms 'galactic cluster' and 'galactic supercluster'?

Extended Task:

You will follow through a European Space Agency developed project to determine the age of the Universe. You will need to learn to manipulate astronomical data using SalsaJ- ESA's version of LTImage.

Lots of the calculations that you are asked to do are beyond the scope of what we are trying to achieve here, so use the spreadsheet to record your data, and the maths will be done for you.

Remember, before you do anything, read, the accompanying information.

To complete the project you will need:

Project instructions and information and the astronomical toolkit

Excel spreadsheet to record and analyse your data

Image data for the 3 cepheid variables you will be analysing.

Bishop pp193-198, Kirk pp129, 135

Back to Astrophysics Option

The Cosmic Microwave Background as Evidence for the Big Bang

Penzias and Wilson discovered that there is a wierd microwave signal coming from whatever direction you look in.

The cosmic background radiation intensity was measured by the Cosmic Background Explorer (COBE) satellite.

This picture was taken in radio wavelengths, and the fact that it shows colour demonstrates that there is an overall 'background' radiation present in the universe. If we consider this radiation as a 'black body', then the radiation profile fits that of deep space having a temperature of 2.7K.

If deep space has a temperature, then it has an energy. But where did this energy come from? Answer the Big Bang.

Current explanation is that based on the solution to Olbers paradox put forward by Hubble, that the Universe is expanding, it must have started in one place, a Singularity. On explosion, the Big Bang, all of the energy (and mass) that exists in the Universe today was released. This implies that the big bang was a very small, very dense, very hot place, and since this time the universe has been getting bigger, less dense and cooler.

What we see now as the Cosmic Microwave Background (CMB) is this energy, spread out over the whole volume of the Universe.

The Hot Big Bang

The fact that the Big Bang was hot and dense is important since explains the large amounts of Helium observed in the universe.

The observed Helium content of the universe is about 27%.

But the amount that could have been produced by hydrogen burning in stars since the big bang could only amount to 2-3%. Accept this by calculations of clever folks!

So at the big bang temperatures and pressures were great enough that a large amount of hydrogen underwent fusion.

We believe that the nuclei for the H and He atoms were formed in the first few minutes after the big bang, but because the temperatures were so high these atoms were completely ionised and thus only existed as nuclei. This is the fourth state of matter - Plasma, an ionised gas.

In this state the universe was opaque, because photons could not travel very far before interacting with electrons (which because of the plasma were everywhere).

After a million years had the universe expanded enough that the temperature had cooled to 3000K. At this point the electrons were captured by nuclei to form gases and there was enough space for the photons in the universe to move freely. Thus making the universe transparent.

Since this time the main events in the universe have been caused by gravity. Fluctuations in the density of regions of the universe have been accentuated by gravitational effects and caused the universe to become 'lumpy'. These lumps on different scales form stars, galaxies and clusters of galaxies (ref: Prof George Smoot).

For lumps to have happened the early universe cannot have been uniform. This is confirmed not only by the existence of density variations throughout the universe but also by variations in the microwave background. COBE detected variations of 3 x 10^-5 K.

So this can help us to make some guesses about the first few minutes of the big bang. Modern particle physics helps us to back up these guesses with evidence.

Tasks:

Watch the video: BBC:Lost Horizons- The Big Bang

Work through the comprehension sheet about the first few minutes after the big bang, answering the two questions

Bishop pp199, 203-211, q5, 6 pp215

Back to Astrophysics Option

The Future of the Universe

The fate of the universe will be decided by gravity and ultimately depends on how much matter there is in the universe. There are 3 possible scenarios.

1. The open or unbound universe. This means that the universe will continue to expand forever. For this to happen there must be so little matter in the universe that gravitational attraction will never have enough effect to stop it expanding.

2. The closed or bound universe. This mean that the universe will stop expanding at some point in the future, and begin to contract again until it ends up at a singularity again à the big crunch. For this to happen there must be more than a certain critical amount of matter in the universe.

3. The flat or marginally bound universe. This means that there is just enough matter in the universe to stop it expanding but not sufficient to start it contracting. If this is the case then it will take forever to stop expanding, but the rate of expansion will slow down as it approaches a definite limit.

To find out the mass of the universe so as to deduce which of the fates awaits us we need to know the average density of the universe. The density needed for the universe to be flat is called the critical density , r_c.

Above the critical density gives us a closed universe, and below it gives us an open universe.

Calculating the critical density and average density of the universe is difficult and involves lots of uncertainty.

Task:

Read: Calculating the critical density of the Universe

Dark Matter/ Dark Energy

Suggestions are that we are missing a lot of the matter, which would be necessary for a closed or flat universe, hence the suspicion of the existence of dark matter. If this is the case, then dark matter/ dark energywould need to make up about 90% of the universe to give us a closed or flat fate.

What is this dark matter (or more recently superseded dark energy)? WIMPs, Brown Dwarfs, etc.

Watch the video ‘The Mystery of Dark Matter'

Task:

Complete worksheets on the video to consolidate your understanding of Dark Matter and DM Candidates: MACHOS, WIMPS, Brown Dwarfs and Black Holes.

Question:

What is your prediction (informed by research) about the fate of the universe: Open, Closed or marginally bound?

Watch the short movie clip of Dr Chris Lintott.

Task:

Read: Recent evidence for an open universe.

Now you're ready for the Universe song!

Bishop pp 199-202

Back to Astrophysics Option

The Relevence and Finance of Astrophysics- Final Tasks

Learning Objectives:

Task: Find one article relating to each of the above aspects of Astrophysics, precee it, and comment.

Back to Astrophysics Option

anrophysics 2007