Cosmology - Life of Stars by Alastair Leith

Online Astronomy Society

Academy

Short Course â&#x20AC;&#x201C; The Life of Stars

By Hugh Allen MA

Contents Online Astronomy Society Academy ............................................................................................................................................................................. 1 Short Course – The Life of Stars .................................................................................................................................................................................... 1 1. Introduction ................................................................................................................................................................................................................. 4 2. The nature of stars ....................................................................................................................................................................................................... 5 3. Energy in – nuclear fusion and nucleosynthesis ....................................................................................................................................................... 12 4. Energy out – starlight ................................................................................................................................................................................................ 14 5. Stable equilibrium – the Main Sequence lifetime ..................................................................................................................................................... 17 6. Classification of stars – OBAFGKM and colour-magnitude diagrams .................................................................................................................... 18 8. The Death of stars ..................................................................................................................................................................................................... 27 9. Supernovae ................................................................................................................................................................................................................ 28 10. Planetary nebulae .................................................................................................................................................................................................... 31 12. References ............................................................................................................................................................................................................... 35

What you will learn in the course -

The basic composition and processes inside stars How mass determines the life of stars How stars influence the appearance of many of the objects visible in amateur telescopes

Links to other OAS Academy short courses -

The Hertzsprung-Russell Diagram Spectroscopy Galaxies: The view from a distance

About the author Hugh Allen is a passionate scientist specialising in chemistry. Growing up in Reading, he attended Reading Grammar School and studied Natural Sciences at Downing College, Cambridge. Hugh has spent more than 25 years in the chemical industry but has always enjoyed a fascination with astronomy and space exploration. Fascination became a passion when his wife bought him a big 8â&#x20AC;? aperture telescope! Hugh is a member of the Online Astronomy Society and the William Herschel Society in Bath. As well as observing and imaging whenever cloud cover permits, he enjoys communicating his passion through writing with an aim to inspire a wider audience and bring astronomy to the forefront. http://www.facebook.com/groups/2217131413/#!/hugh.astrophotography

Online Astronomy Society Academy Short Course – The Life of Stars

1. Introduction No telescope is needed to see stars on a clear, dark night. Four thousand years ago the ancient Egyptians believed they were lanterns hanging from the roof of the Universe. The speculations of ancient Greek philosophers from 2,500 years ago included pinpricks in a canvas that shielded us from the celestial fire, or pinheads attached to a rotating celestial sphere; the earth-centred celestial spheres of Pythagoras, Aristotle and Ptolemy became deeply embedded in European Medieval thought. It took a 16th century Polish priest called Copernicus to begin the scientific revolution by placing the Sun at the centre of our solar system. Yet it was another three hundred years before the similarities between the stars and our own Sun were conclusively demonstrated. Mapping the stars has always been an important human endeavour, to measure the passage of time and to aid navigation. It is perhaps no surprise that the patterns (constellations) made by the stars would become the landmarks on such maps. The naming and boundaries of the 88 modern constellations are now controlled by the International Astronomical Union (IAU). They are substantially based on their predecessors from Roman and Greek times which were themselves derived from the work of Babylonian astronomers around 600BC, in what is now Iraq. Figure 1 shows a map of some of the modern constellations. Their Latin names are derived from Greek mythology and from the shapes that they represent. Astronomy’s long heritage is also found in the naming of the brightest stars on the map. The Arabic origin of many names is betrayed by the prefix ‘Al’ e.g. Alnitak (the girdle) in Orion, Aldebaran (the follower) in Taurus, Alrescha (the well-rope) in Pisces and Almach (the desert lynx) in Triangulum. Arabic astronomy flourished between the 8th and 15th centuries and left a rich legacy. We now estimate that our galaxy the Milky Way contains around 100 billion stars. And the Universe as a whole contains an estimated 2 x 1021 stars (that’s two thousand billion billion, if that helps comprehend such an enormous number). The life of stars is hugely varied and complex and this course aims only to present some of the key, underlying themes. And rest assured, the life of stars is a visual as well as a scientific treat!

Figure 1:

A view of some of the modern constellations in the northern hemisphere at 22:00 GMT on 1st January 2013, looking south from latitude 51째 North

2. The nature of stars To understand the nature of stars we need to know their mass. The mass of our Sun is fairly modest at just under 2 x 1030 kg (or 2,000 billion billion billion tonnes). The Sun makes a very useful reference point for understanding other stars so the Sun’s mass is conventionally given its own symbol Mʘ. Relative to the Sun, the mass of other stars shows a wide variation between 0.08Mʘ up to about 150Mʘ. The frequency at which they are observed in our neighbourhood of the Milky Way has an interesting distribution: Mass range relative to the Sun’s mass, Mʘ 10 - 150 2 – 10 0.5 – 2 < 0.5

Approximate relative frequency 1 10 50 Several hundred

Much of the mass of a star is concentrated in its hot, dense core. The Sun’s core occupies just 10% of the volume (or ~25% of the 700,000km radius Rʘ) but it contains about 40% of the mass. The temperature of the core is 15 million Kelvin (0 Kelvin = minus 273°C = absolute zero) whilst the outer surface of the Sun is at ‘only’ 5,777 Kelvin. The composition of the Sun’s mass is also interesting to think about. Here are the Sun’s ‘top ten’ chemical elements in order of mass(1) : Element Hydrogen Helium Oxygen Carbon Nitrogen Silicon Magnesium Neon Iron Sulphur

Number of protons in the nucleus 1 2 8 6 7 14 12 10 26 16

Number of neutrons in the nucleus 0 2 8 6 7 14 12 10 30 16

% by mass 71.0 27.1 0.97 0.40 0.096 0.099 0.076 0.058 0.014 0.004

% by number of atoms 91.2 8.7 0.078 0.043 0.0088 0.0045 0.0038 0.0035 0.030 0.015

The Big Bang model predicts that at the birth of the Universe hydrogen and helium would have been formed in a mass ratio of 3:1, along with traces only of deuterium (one proton and one neutron) and lithium (3 protons and 4 neutrons). No other chemical elements 6

were formed at the Big Bang and there is observational evidence that supports these predictions (2). Chemical elements heavier than hydrogen and helium were already part of the material from which the Sun was born.

By taking a moment to gaze at the night sky it is clear that stars have widely differing brightnesses and subtly different colours. The subtle, contrasting tints of stars are best picked out when they are near to each other in the sky. For example there are two bright stars which frame the unmistakable constellation Orion: orange-red Betelgeuse above and blue-white Rigel below. Here is how my Meade DSI II colour camera picks them out through my 8â&#x20AC;? Meade LX90 reflector telescope (all of the images in the course were captured with this equipment):

Betelgeuse - 640 light years distant Apparent magnitude 0.5 Surface temperature = 3,500K

Rigel - 860 light years distant Apparent magnitude 0.12 Surface temperature = 11,000K

Note that these stars are too far away to be seen as a disc, they are actually just points of light. The images appear disc-like only because the pixels in the camera have become saturated by the brightness of the stars. At a given exposure setting, a more faint star produces a smaller disc.

The system of apparent magnitudes to compare the brightness of stars was developed around 2,000 years ago by Hipparchos in ancient Greece. The faintest naked eye stars were put in the 6 th class whilst Sirius, the brightest star, was put in the 1st class. This ancient system of classification has now been transformed into a quantitative measurement system. The bright star Vega is the reference standard with a magnitude of zero (almost). Each step in magnitude away from Vega is equivalent either to a 2.512 reduction in brightness (for +ive magnitudes) or increase in brightness (for -ive magnitudes). A difference of 5 steps on the magnitude scale is therefore equivalent to a 100-fold difference in brightness (2.5125). The brightness of stars that we observe in the sky is described as the apparent magnitude (m) because it does not take into account the distance to the star. The brightness of a star falls away with the square of its distance from us (the inverse square law). And this brings us to some fascinating history about the measurement of the distance to stars using parallax anglesâ&#x20AC;Ś.. There is no need to understand the maths to appreciate the power of parallax. Just hold a finger in front of your eyes and watch it move against the distant background as you open one eye and shut the other. The angle through which it moves is the parallax angle. Now move the finger much closer to your eyes and see how much more it moves against the background. Imagine that your eyes are replaced by the position of the Earth on opposite sides of its orbit around the Sun. This is the baseline for parallax measurements in 8

astronomy and the geometry was understood as far back as the ancient Greeks. The difficulty is measuring the angle; the distances to even the nearest stars are so enormous that the parallax angles are tiny, but the scientific rewards are enormous too!

The race to measure the first reliable parallax angles was a fascinating episode in the history of astronomy(3). It was eventually won by the German astronomer Friedrich Bessel in 1838 who measured a parallax angle of just over 0.3 arcseconds for the star 61 Cygni (see left). This is equivalent to a distance of about 10.4 light years. (1 arcsecond is 1/60th of 1/60th of a degree!). Our nearest neighbour Proxima Centauri produces a parallax angle of only 0.7687arcseconds which puts it at the huge distance of 4.24 light years (light travels about 9,460 billion km in a year). The accuracy of parallax angle measurements is best improved by getting above the interference of the Earth’s atmosphere allowing measurement of more distant stars. The reach of parallax measurement was extended to a distance of 1500 light years in 1989 with the launch of the Hipparcos satellite telescope, achieving an accuracy of milliarcseconds. The European Space Agency’s Gaia mission launching in October 2013 will have an accuracy of microarcseconds, allowing parallax measurements to tens of thousands of light years distance.(4) 61 Cygni (Piazzi’s Flying Star) – a beautiful binary star

Knowing the distance to a star it is possible to calculate its absolute magnitude, M, defined as if the star was placed at a reference distance of 10 parsecs or 32.6 light years (1 parsec is defined as the distance to an object which would produce a parallax angle of 1 arcsecond). The absolute magnitude allows the brightness of stars to be compared directly to each other. And once we know the star’s absolute magnitude it is possible to calculate its luminosity, L. This is the total amount of light energy/second emitted by the star. The luminosity of the Sun is about 3.9 x 1026 Watts and is given the symbol Lʘ. Like solar mass Mʘ discussed at the start of this section, Lʘ makes a useful reference to help compare different stars. Here is a comparison of some stars in order of their mass relative to Mʘ. All are visible to the naked eye except our nearest neighbour Proxima Centauri: Star name Proxima Centauri The Sun Procyon Altair Sirius Vega Alnair Achernar Sigma Orionis A

Mass Luminosity L Surface (relative to Mʘ) (relative to Lʘ) temperature, K 0.12 0.0017 3,042 1 1 5,777 1.5 7.7 6,530 1.8 10.6 7,550 2.0 25 9,940 2.1 37 9,600 4 380 13,500 8 3300 20,000 18 35000 32,000 10

Radius, R (relative to Rʘ) 0.14 1 2 1.8 1.7 2.3 3.6 10 ?

Distance (light years) 4.24 0.000016 11.4 16.8 8.6 25.3 101 144 1150

Apparent magnitude, m +11.05 - 26.74 + 0.4 + 0.76 - 1.44 + 0.03 + 1.73 + 0.45 + 4.20

Absolute magnitude, M + 15.5 + 4.83 + 2.68 + 2.20 + 1.45 + 0.58 - 0.73 - 2.77 - 5.25

By knowing more about their nature you can begin to see the emergence of some order to the life of stars e.g. combining this data with the table at the start of the section shows that most of the mass is found in small stars, but most of the light is from giants.

3. Energy in – nuclear fusion and nucleosynthesis Stars are powered by the energy released from nuclear fusion reactions in their hot, dense cores. Small nuclei fuse together to form heavier chemical elements in a process called nucleosynthesis. The products of nuclear fusion have a lower nuclear binding energy than the starting materials, detected as a lower total mass. The missing mass, known as the mass defect, has been converted to energy according to Einstein’s famous equation: energy,E = mass,m x square of the speed of light,c². The energy is released principally as photons of electromagnetic radiation called gamma rays, the most energetic form of light. The principal fuel for these fusion reactions is hydrogen, the most abundant chemical element in the Universe. All of the heavier elements that are now found in the Universe are the product of nucleosynthesis inside stars. Astronomers refer to all of these heavier elements as metals, although my chemistry education makes it difficult to write that phrase! The temperature in the core of the Sun is around 15 million Kelvin. It is maintained by the energy released from the fusion of hydrogen (one proton in the nucleus) to form helium (two protons and two neutrons) in the proton-proton cycle. There is a strong repulsive force between two positively charged protons, known as coulombic repulsion. But at temperatures above 10 million Kelvin protons have enough kinetic energy to overcome coulombic repulsion when they collide. The proton-proton cycle has several steps but the overall process is: 4x hydrogen nuclei (protons) → 1x helium nucleus (2 protons, 2 neutrons) + 2x neutrinos + energetic photons (gamma rays)

At core temperatures above 15 million Kelvin the fusion of protons can also be catalysed by the presence of atoms of carbon, nitrogen and oxygen in what is known as the CNO cycle. An evocative way to think of both of these processes is that hydrogen is ‘burnt’ to leave an unreactive helium ash. In every second, our Sun fuses nearly 614 million tonnes of hydrogen into 610 million tonnes of helium (remember the mass defect, and of course the much smaller mass lost as neutrinos). If the temperature increases further then a complex web of nucleosynthesis reactions becomes possible. One of the most significant is the triple-alpha process which fuses three helium nuclei into carbon (6 protons and 6 neutrons). The triple-alpha process requires a core temperature of at least 100 million Kelvin. At even higher temperatures carbon can be converted into heavier elements by the alpha process. Successive fusion with helium nuclei produces first oxygen, then neon, magnesium, silicon, sulphur, argon, calcium, titanium, chromium, and finally iron. Why finally? Because iron has the lowest nuclear binding energy of any chemical element. For nuclei heavier than iron the binding energy increases which means energy must be absorbed not released in their formation. Nucleosynthesis in the core of stars essentially comes to a halt with the formation of iron. As we shall see later, elements heavier than iron are formed by absorbing some of the huge quantities of energy released in violent stellar processes such as supernovae. This is known as explosive nucleosynthesis. As they travel outwards from the core the gamma rays produced in the fusion process are absorbed and re-emitted multiple times by atoms in the star. The energy of the re-emitted photons is directly related to the temperature. Since the temperature of the star falls dramatically towards its surface so therefore does the individual energy of the re-emitted photons (but remember energy is always conserved so the number of re-emitted photons is higher). Many thousands of years after its release in the core as gamma rays, the 12

electromagnetic energy eventually reaches the visible outer layer of the star called the photosphere and exits into outer space. By this point much of the energy is concentrated into the UV-visible-infrared part of the electromagnetic spectrum.

4. Energy out â&#x20AC;&#x201C; starlight The energy from fusion reactions in the core of a star is nearly all released as photons of light through the photosphere. The rate at which the energy is released is known as the luminosity of the star and is measured in Watts. The luminosity and its spectral distribution can be modeled by the physics of black bodies i.e. they both depend on the temperature of the photosphere. The higher the surface temperature the greater the luminosity and the greater the proportion of energy in higher energy photons (shorter wavelengths). These relationships can be represented graphically: The electromagnetic spectrum

Black body radiation curves for specific temperatures (the area under the curves represents the luminosity)

Image Credit: http://en.wikipedia.org/wiki/File:Electromagnetic-Spectrum.png

Take another look at the surface temperature information with the images of Betelgeuse and Rigel in Section 2. The subtle tints of stars are directly related to their surface temperatures. With a surface temperature Tʘ of 5,777 Kelvin the Sun’s emission intensity peaks at a wavelength of 502nm, represented by the symbol λmax ʘ. The peak emission wavelength λmax of other stars can then be calculated using the simple relationship (5) : λmax = λmax ʘ x Tʘ/T where T is the surface temperature of the star The subtle colours of the stars pictured in Section 2 can now be explained. For Betelgeuse with a surface temperature of 3,500 K the emission intensity peaks at 829nm in the near infrared, which gives the star an orange-red tint to the naked eye. And for Rigel with a surface temperature of 11,000 K the peak is at 263nm in the UV part of the spectrum, hence the blue tint.

The visibility of sunspots on the Sun is a wonderful manifestation of the relationship between temperature and luminosity. When we look at the Sun (not directly of course!) we are looking at the photosphere. Sunspots are cooler zones in the photosphere which appear dark because they have lower luminosity in contrast to their surroundings.

The Sun at 14:00 GMT on 13th January 2013. Active Region AR 11654 is very prominent, left of centre of the image

The chemical composition and temperature of the photosphere lead to the presence of characteristic absorption lines in the light spectrum. Spectroscopy is used to study the spectrum of stars and is a vast and fascinating topic in its own right. It is treated in detail in a separate OAS Academy course. The luminosity, surface temperature and radius of a star are determined by its mass. It is interesting to look at the equations which link them. Their simplest forms use values relative to those of the Sun. The following equations are reasonable approximations only for stars in the mass range 0.8Mʘ to 10Mʘ but they give a good feel for how fundamental is the mass of the star: L / Lʘ ≈ (M / Mʘ)3.5

=> a star with twice the mass of the Sun will have approximately 11 times the luminosity 15

(T / Tʘ)4 ≈ (M / Mʘ)2.5

=> a star with twice the mass of the Sun will have approximately 1.5 times the surface temperature

R / Rʘ ≈ M/ Mʘ

=> a star with twice the mass of the Sun will have approximately twice the radius

5. Stable equilibrium – the Main Sequence lifetime For the bulk of their lives stars are in hydrostatic equilibrium. The force of gravity from the mass of the star pulls inwards causing the star to collapse and heat up. The energy from the fusion reactions in the core produces an outward thermal pressure causing the star to expand and cool. A star is stable because the competing forces of gravity and thermal pressure are in equilibrium. The equilibrium is self-regulating because the rate of the fusion reactions is very sensitive to the core temperature CT; the rate of the protonproton cycle varies approximately with CT4 whilst the CNO cycle varies with CT15 !! When a star is in a state of hydrostatic equilibrium it is said to be on the Main Sequence. The position of the equilibrium depends again on the mass of the star. Heavier stars require a higher outward thermal pressure to resist the increased force of gravity and must therefore burn their nuclear fuel at a faster rate. This hints at a surprising outcome for the comparative lifetime of stars on the Main Sequence. The nuclear fuel available to support the equilibrium is proportional to the mass of the star, and the mass of fuel is proportional to the energy which it can produce (E = mc²) which is all released as luminosity. This leads to the following relationship between the Main Sequence lifetime and the mass, relative to the Sun: Main sequence lifetime, t ≈ Main sequence lifetime, tʘ x (Mʘ/ M)2.5 A surprising consequence is that a more massive star will have a considerably shorter lifetime. The estimated Main Sequence lifetime of the Sun tʘ is 10 billion years (1 x 1010 years) so we can predict the following lifetimes for other stars of different masses: Mass M relative to Approximate Main the Sun’s mass Mʘ Sequence lifetime (years) 0.08 5.5 x 1012 0.5 5.7 x 1010 1 1.0 x 1010 2 1.8 x 109 10 3.1 x 107 150 3.6 x 104 The smallest stars can expect to be on the Main Sequence for trillions of years (one trillion is 1 x 10 12) which is a hundred times longer than the current age of the Universe (13.7 billion years or 1.37 x 10 10). The largest stars will live for less than 10 million years (1 x 107 years) which is the cosmic blink of an eye.

6. Classification of stars – OBAFGKM and colour-magnitude diagrams Since 1943, stars have been classified by their colour (or surface temperature) according to the Morgan – Keenan system. Each spectral class is broken into 10 subdivisions 0 – 9 going from hotter to cooler: Spectral class

Temperature range (Kelvin, K)

50,000 – 28,000

28,000 – 10,000

10,000 – 7,500

7,500 – 6,000

6,000 – 4,900

4,900 – 3,500

3,500 – 2,000

Colour description

Blue

Blue - white

White

White-yellow

Yellow

Orange

Red

Example

Theta Orionis C Rigel Vega Polaris, The Sun Arcturus Betelgeuse (Class O6) (Class B8) (Class A0) (Class F7) (Class G2) (Class K1) (Class M2) in the Orion Nebula ‘Oh Be A Fine Girl Kiss Me’ is often given as a useful mnemonic to remember the different classes. Can you think of a different one?

Going a step further, the colour of a star can be quantified by relating its magnitude in a blue filter B (peak transmission at 440nm wavelength) to its magnitude in a yellowgreen filter V (peak transmission at 550nm which corresponds to the human eye’s peak visual response). The star’s colour index (C.I.) is calculated as B – V. A blue star has a more negative C.I. then a red star. A star of spectral class A0 has a C.I. of 0.0. This allows us to plot what is known as a colour-magnitude diagram. I created this one myself from publicly available Hubble Telescope B/V magnitude data for stars in the open cluster NGC 2266. In an open cluster the effect of distance can be ignored since all the stars are clustered together at a similar distance, about 11,000 light years in the case of NGC 2266. The y-axis is the magnitude of the stars measured through the V filter but plotted in reverse so that brighter stars are more naturally higher up. The x-axis is the Colour Index. Most of the stars are concentrated around a straight line from top left down to bottom right. Stars on or near this line are on the Main Sequence – their brightness and surface temperature are determined by their mass. There are no stars on the line with a C.I. lower than about +0.1. This is how we know the cluster is ~1 billion (1 x 109) years old. The stars not on the Main Sequence are also not randomly distributed. Their position tells us something about how stars evolve at the end of their lives. The importance of the relationship between magnitude and colour in 18

understanding the life of stars was first recognised by Ejnar Hertzsprung and Henry Russell in 1910. The graph is also known as a Hertzsprung-Russell diagram and such diagrams are discussed in more detail in another OAS course.

7. A cold, dark birthplace Stars are born inside clouds of very cold gas and dust. The clouds can fill huge volumes of space hundreds of light years across such as the Orion Molecular Cloud. Others are much smaller. Here is an attempt to classify the different scales which can be observed: Mean diameter (light years) Density, number of hydrogen molecules/cm3 Mass, relative to Mʘ Temperature, Kelvin

Giant molecular cloud 130 100

Molecular cloud 32 300

Molecular clump 13 1000

Cloud core 0.5 100,000

100,000 15

10,000 10

1000 10

10 10

Data from P.H. Bodenheimer, Principles of Star Formation, Astronomy and Astrophysics Library, DOI 10.1007/978-3-642-15063-0 2, © Springer-Verlag Berlin Heidelberg 2011

Compared to the gas densities we are used to here on Earth (~2.5 x 1019 molecules/cm3) these clouds are very tenuous! The clouds’ extent can be mapped from the rotational microwave spectrum of carbon monoxide (CO) which is present in the clouds at trace levels. Dust in the clouds is made of carbon and silicate materials and is extremely fine at around 0.1µ in size. If the cloud is sufficiently deep then the dust can obscure the visible starlight behind it making the silhouette of the cloud visible even in amateur telescopes, described as a dark nebula. The late 19th century astronomer E.E.Barnard made a photographic survey of the Milky Way from which 370 dark nebulae were identified and catalogued. To the left is my own image of the 139th dark nebula in Barnard’s catalogue. The angular length of the cloud is ~20 arcminutes. I could not find any reference to the distance to the cloud, but where distances to Barnard objects are referenced they seem to lie between 400 and 1500 light years. So let’s assume that Barnard 139 lies at about 1000 light years. That would make the cloud around 6 light years long.

Barnard 139 by Hugh Allen

At the exceedingly low temperatures in these clouds there is an equilibrium between the weak force of gravity pulling the gas and dust together and the outward thermal pressure of the material due to its temperature. But the equilibrium can be disturbed in favour of

gravitational attraction when a local increase in density occurs, caused perhaps by clouds colliding or by the shock wave from a nearby supernova explosion. The disturbed region of the cloud begins to collapse on itself – star formation has begun. As the cloud collapses the gravitational potential energy is converted into heat and the collapse continues until the thermal pressure can balance the collapse. The temperature needed to balance the gravitational collapse will depend on the mass of the collapsing material. If it becomes hot enough it will begin to emit at optical wavelengths to form a protostar (black body radiation). If the temperature within the protostar eventually reaches ~10 million Kelvin nuclear fusion of hydrogen begins. The energy released by the fusion process is enough to counterbalance gravity and the system returns to equilibrium. A new star is born, stable for as long as its nuclear fuel can burn. A protostar with less than 0.08 solar masses never reaches the 10 million K core temperature needed for hydrogen fusion. There is not enough gravitational energy from the collapse to heat the core sufficiently. The result is a “failed star” called a brown dwarf whose surface is so cool that it radiates light mainly in the infrared. As infrared astronomy has developed several new spectral classes have been added to OBAFGKM. There may be many more of these ‘failed’ stars than any other type: L class: Surface temperature 1300 – 2000 Kelvin T class: Surface temperature 700 – 1300 Kelvin Y class: Surface temperature <600 Kelvin The estimated time for the entire process of initial cloud collapse to the initiation of nuclear fusion is strongly dependent on the mass of the newborn star: Mass of newborn star relative to the Sun, Mʘ 0.5 1.0 3.0

Time from onset of collapse to initiation of nuclear fusion 160 million years 50 million years 2½ million years

As the cloud collapses it starts to spin and flatten forming a disc around the protostar. This is the origin of planet formation. Soon after the birth of new stars they are still surrounded by the gas and dust from which they were born producing spectacular nebulae. The intense light from the largest newborn stars is scattered by the dust creating a beautiful blue colour. Hydrogen gas is ionized by UV light producing characteristic red hydrogen-alpha emission at a wavelength of 656nm. Where the surrounding dust is optically thick it produces beautiful clouds and filaments. The nebulae will fade as the stellar winds from the newborn stars drive away the remaining gas and dust. One of the best known star forming regions visible to the naked eye is the Orion Nebula (Messier object M42), the middle ‘star’ of Orion’s sword. All are beautiful in a telescope, especially in photographic images:

Orion Nebula M42

Triffid Nebula M20

Iris Nebula NGC 7023

The volume of collapsing material may be small resulting in the birth of a single star. More commonly, stars form within larger clouds in groups or clusters as the collapsing cloud fragments. Evidence for this is all around us. Approximately 50% of stars are actually part of a double (binary) or multiple star system, where the stars are orbiting around a common centre of gravity. The frequency of multiple star systems increases with stellar mass. In lower mass stars around 10 â&#x20AC;&#x201C; 30% are in binary or multiple star systems, increasing to close to 100% for the most massive stars. Here are some beautiful examples:

40 Eridani a triple star system only 19.5 light years away

Albireo one of the prettiest colour-contrast binaries in the sky

The Pole Star, Polaris

The sky is also sprinkled with many clusters of young stars, described as open clusters. One of the most famous is the Pleiades or Seven Sisters, visible to the naked eye on dark winter evenings. These clusters are only loosely bound by gravity and will eventually drift apart in their tour around the galaxy. It is believed that the Sun was born inside such a cluster nearly 5 billion years ago (6) :

Open cluster M52

Open cluster NGC 2244

The Trapezium Cluster lies at the bright core of the Orion Nebula

8. The Death of stars Unsurprisingly, the fate of a star is determined by its mass. Here is a simplified view of the outcomes: Small: 0.08 - 0.5Mʘ Main sequence Fusing H to He

Intermediate: 0.5 - 8Mʘ Main sequence Fusing H to He.

Massive: >8 Mʘ Main sequence Fusing H to He

Red giant phase H fuel in the core is exhausted Core contracts H fusion re-starts in a shell around the core Envelope cools and expands to ~100 times main sequence radius Significant mass loss from the surface as a solar wind (maybe 30%)

Red giant phase H fuel in the core is exhausted Core contracts He fusion to carbon begins inside H shell He fuel in the core is exhausted Core contracts Carbon fusion to oxygen begins inside He shell Carbon fuel in the core is exhausted Core contracts Oxygen fusion to neon begins inside carbon shell etc. etc

He core reaches 100 million K He fusion into carbon and oxygen begins (the helium flash) Envelope contracts and heats up

Eventually iron builds up in the core and fusion stops Iron core can no longer resist massive gravitational collapse Rebound causes a gigantic explosion releasing huge amounts of energy and leaving an extremely dense remnant core

Red supergiant phase He fuel in the core exhausted He and H fusion re-starts in shells around the core. Envelope expands and cools again Pulsations lead to further cycles of mass loss

Remnant He white dwarf H exhausted Star too cool to initiate He fusion

Core collapse supernova, Type II

Original stellar mass < ~50Mʘ Remnant neutron star Core collapse compresses protons and electrons into neutrons and neutrinos One sugar cube-sized piece of neutron star weighs ~500 billion kg The radio emissions from neutron stars can be observed from Earth as pulsars

Remnant carbon-oxygen white dwarf Extremely hot, dense carbon-oxygen core at ~100,000K, left exposed to slowly cool Remaining envelope ejected into space Nearby ejected material emits visible light, due to atomic excitation by intense UV light from hot core = planetary nebula If in a binary system, the dwarf can accrete mass from its companion, up to 1.4Mʘ = Chandrasekhar limit = Supernova, Type 1a no significant remnant

> ~50Mʘ Remnant black hole Collapse compresses neutrons into an infinitely dense singularity Black holes can be observed from the electromagnetic radiation emitted by matter as it spirals towards the black hole

9. Supernovae The initial ejection velocity from a supernova is around 10,000 km/sec. Within a few days a supernova can outshine an entire galaxy of stars. The huge amount of energy released by the collapse leads to explosive nucleosynthesis and the formation of chemical elements heavier than iron including many radioactive nuclei. The release of energy by the supernova is prolonged by the decay of these radioactive nuclei, particularly radioactive nickel and cobalt. In our Milky Way galaxy it is estimated that there should be on average one supernova every 100 years. Unfortunately, since Kepler’s supernova of 1604 none have been visible from our earthly vantage point. Supernovae are however regularly observed in other galaxies. In 2011 there were two relatively bright explosions, in the Whirlpool Galaxy M51 and the Pinwheel Galaxy M101. I was able to make ‘before’ and ‘after’ image comparisons which clearly show the supernovae’s locations. On the left is my comparison image of the supernova from June 2011 in M51, at a distance of about 34 million light years. I caught the supernova almost at the moment of peak brightness. Type II supernova in the Whirlpool Galaxy M51

Early in 2012 I plotted the light curves of each supernova to see how they had evolved. Apparent visual magnitude data is publicly available on the website of the American Association of Variable Star Observers (AAVSO). Supernova light curves have characteristic shapes and peak intensities which are typical of the type of precursor star (note that these curves do not take into account the different distances to the galaxies). SN2011fe in M101 was a Type 1a supernova caused by mass accretion onto a white dwarf from a nearby companion star. The supernova peaked at an apparent magnitude of about 10.0, not far behind the 8.3 apparent magnitude of the galaxy as a whole. The precursor to SN2011dh was a massive star leading to a Type II core collapse supernova. I captured my image of sn2011dh as it was still 28

climbing towards its peak brightness. Yet even at a distance of ~23 million light years it appears at least as bright as any of the foreground curtain of nearby stars in the Milky Way. At least 90% of the mass of the star is ejected in a supernova explosion (the so-called mass cut). The expanding shell of debris from a supernova remains visible for tens of thousands of years. There are some famous examples that can be observed in amateur telescopes e.g. the Crab Nebula M1 and the Veil Nebula. The Veil Nebula spans nearly 3 degrees of sky (the Full Moon is only about half a degree in diameter) which at a distance of ~1400 light years means that the shell of debris is now ~75 light years wide! The Type II supernova is estimated to have exploded some 8,000 years ago and the shell is still expanding at around 170 km/sec. I recently captured images of two of the brightest parts of the shell:

The bright disc in the right hand photo is the foreground star 52 Cygni, overexposed to bring out the faint nebula

10. Planetary nebulae Planetary nebulae are amongst the most beautiful but ethereal objects in the sky. They are too faint to be seen with the naked eye but in a telescope many become clearly visible, at least under dark skies. Enduring for just a few tens of thousands of years planetary nebulae represent the final stage in the life of intermediate mass stars like our own Sun. They were first observed in the 18th century and the term was coined by William Herschel in 1784 because of their disc-like appearance in the telescopes of the time. Their true nature was not conclusively proven until the following century when spectroscopy was developed. An early pioneer of astronomical spectroscopy was the English amateur astronomer William Huggins. During the 1860’s the field of astronomical spectroscopy was in its infancy after the formulation of Kirchoff’s Laws in 1859. These laws described the conditions in which chemical elements produce unique finger prints of matching absorption or emission lines in their light spectrum. The vast potential for analysing the light from astronomical objects was immediately obvious and Huggins was one of several pioneers, supported by his next-door-neighbour the King’s College Professor of Chemistry W. A. Miller. I’ll let Huggins himself describe his first observation of the spectrum of a planetary nebula: “My surprise was very great, on looking into the small telescope of the spectral apparatus, to perceive that there was no appearance of a band of coloured light such as a star would give, but in place of this, there were three isolated bright lines only. This observation was sufficient to solve the long agitated inquiry in reference to this object at least and to show that it was not a group of stars but a true nebula….[composed] from glowing or luminous gas’. This is an image of the small planetary nebula NGC 6804. At a distance of about 4,200 light years and with an angular diameter of just over 1 arcminute (1/60th of a degree) it is 1.2 light years in diameter, very typical for a planetary nebula. The dying star is clearly visible at its centre. Below my image I have copied a black & white photo of the nebula’s visible light spectrum taken through a long-slit spectroscope, to which I have added a few annotations. The spectrum is from a study by Justin Cantrell of Georgia State University using the 100” Isaac Newton Telescope on La Palma, Canary Islands.(7) (this telescope was until 1979 sited at Herstmonceux Castle in Sussex, UK). Although hydrogen is the most abundant element in the nebula it is oxygen that produces the strongest emission line at blue-green wavelengths of 495.9nm and 500.7nm. The more faint hydrogen-alpha emission line is at the red wavelength of 656.3nm. 31

An intermediate mass star will have expelled 50% or more of its mass as gas and dust by the end of the planetary nebula phase. The material is ejected from the dying star at a velocity of 20 – 30 km/sec. It can form complex 3-dimensional shapes which are hard to discern. To the left is an image of the planetary nebula NGC 6781. Unusually amongst planetary nebulae it seems to have an almost perfect bubble shape. In fact the nebula is believed to form an open-ended cylinder with a slight hour-glass shape. We happen to be looking down the open end with the cylinder tilted at an angle of just 25° from our line of sight. In the corner I have pasted in the same orientation a model of NGC 6781 from a 2006 paper by Schwarz and Monteiro.(8) Like the faint blue star just visible at the centre of NGC 6781, our own Sun will briefly form a planetary nebula at the end of its life in 5 billion years’ time. Life on Earth will not have survived this final phase of the Sun’s existence.

Planetary nebulae can be observed at different stages in their evolution:

Planetary nebula NGC7027 – less than 1000 years old Some structure just becoming visible Distance ~3,000 light years/angular size ~14 arcseconds => diameter of ~0.2 light years

Planetary nebula M57 the Ring Nebula – estimated age 7,000 years Distance ~2,300 light years/angular size ~90 arcseconds => diameter of ~1 light year

11. Conclusion – the life cycle of stars The life of stars can be portrayed in a sequence of images. You may have realised how the stages in the life of stars have come full circle: Birth inside a cold, dark cloud of gas and dust, a long stable life, death by expelling gas and dust back into the galaxy which can form new stars. I hope you have enjoyed the journey……

17Mʘ O-type star AE Aurigae Distance ~1,500 light years

Star forming nebula of gas and dust Lagoon Nebula M8 Distance ~5,000 light years

Open cluster NGC 2266 Distance ~5,200 light years

1.7Mʘ F-type star Mu Bootis in triple star system Distance 120 light years

Type II supernovae in the Whirlpool Galaxy Distance ~23 million light years Planetary Nebula Dumbbell Nebula M27 Distance ~1,400 light years

34 Supernova remnant Crab Nebula M1 Distance ~6,500 light years

12. References (1) The Sun’s top 10 chemical elements in NASA’s Imagine the Universe http://imagine.gsfc.nasa.gov/docs/ask_astro/answers/961112a.html (2) ‘Astronomers Find Clouds of Primordial Gas from the Early Universe, Just Moments After Big Bang’ Science Daily 11 th Nov 2011 http://www.sciencedaily.com/releases/2011/11/111110142050.htm (3) Alan W. Hirshfield (2002) ‘Parallax: The Race to Measure the Cosmos’ (4) The Hipparcos Space Astrometry Mission http://www.rssd.esa.int/index.php?project=HIPPARCOS&page=index (5) Wavelength intensity applet http://highered.mcgrawhill.com/olcweb/cgi/pluginpop.cgi?it=swf::535::365::/sites/dl/free/0073512176/220727/Blackbody_Nav.swf::Blackbody%20Radiation%20 Interactive (6) Discussion of the Sun’s birth cluster http://www.skyandtelescope.com/skytel/beyondthepage/Mark-Giampapa-on-the-Sun-and-SolarTwins-137116423.html (7) NGC 6804 spectral analysis http://www.chara.gsu.edu/~cantrell/pne/index.htm (8) NGC 6782 shape study Schwarz and Monteiro, ‘3D photoionisation structure and distances of planetary nebulae III NGC 6781’ The Astrophysical Journal vol 648 p430-434. http://www.uninove.br/PDFs/Publicacoes/exacta/exactav4n2/exactav4n2_3e13.pdf