Stellar Evolution : The Life and Death of Our Luminous Neighbors by Arthur Holland and Mark Williams Introduction and Why You Should Care Observations of the stars with calculations of stellar models have give astronomers a comprehension of stellar evolution. Stellar evolution is the process in which the forces of pressure (gravity) alter the star. With these forces acting upon stars, their characteristics change dramatically over the period of their existence. Stellar evolution is inevitable as stars deplete their initial fuel sources. The search for new fuel sources affects the properties of stars as they evolve. This evolution is a process that consists of many different stages with fuel consumption as the dominant life cycles of an evolving star. Stellar evolution, in the form of these fuel consumption stages and their finality, is important because it is responsible for the production of most of the elements (all elements after H and He). Moreover, stages in the life cycle of stars are a vital part in the formation of galaxies, new stars and planetary systems. Time scales of Stellar Fuel Consumption The time scales of stellar evolution depend on the mass of the star. The rule governing stellar evolution is the more mass present, the faster the evolution for the star through the fuel consumption stages. Another property directly linked to the mass and evolution of a star is its luminosity. The mass-luminosity relation demonstrates that the main sequence on an H-R diagram (a chart plotting the the luminosities of stars against their surface temperatures) is a progression in mass as well as in luminosity and surface temperature (Kaufmann,1991, pg. 356). The Hertzsprung- Russell diagram is two possible plots of spectral types dependent on absolute magnitude when compared to temperature and luminosity. These H-R diagrams allow astronomers to conceptualize visually the mass-luminosity relationship as it pertains to the fuel consumption evolution of stars. The H-R diagram allows astronomers to plot visually the different stages in the evolving life cycle of stars. A H-R diagram (the red line represents the Main Sequence Stars). Stages in the Life of an Evolving Star Welcome to the life of a new born star. Main sequence stars are a range of stars based on size and surface temperature starting from the hot, bright, bluish stars in the upper left corner of a H-R diagram to the cool , dim, reddish stars in teh lower right corner of the diagram. Life for new stars begins in the Main Sequence. These mature stars undergo a remarkable transformation after they consume all the hydrogen in their core. With the hydrogen consumed, stars leave the main sequence and expand to form red giants. With this new stage, the fusion of helium begins to form heavier elements like Oxygen and Carbon. This process of expansion- collapse-expansion of stars forms the light elements present in the universe (up to Fe). Life in the Suburbs : a Main Sequence Star. Main sequence stars are stars who�s luminosity and surface temperature place it in the �main sequence� of a Hertzsprung- Russell diagram. A fundamental property of all main sequence stars is thermal equilibrium. Thermal equilibrium is the liberation of energy from the interior of the star balanced by the energy released at the surface of the star. The energy released by a main sequence star is produced by hydrogen burning in its core (the fusion of 4H into 4He). Another fundamental property of a main sequence star evolution is hydrostatic equilibrium. Hydrostatic equilibrium reflects the required pressure in the core of a star to support the weight of the outer plasma layers. The heat produced from hydrogen in the core burning supports this outward pressure upon the outer plasma layers. As a main sequence star depletes the supply of hydrogen in the core, thermal equilibrium unbalances and the pressure in the star�s core lessens. Thermal equilibrium unbalances because the fusion of four hydrogen atoms into one helium atom decreases the number of particles present in the star�s core. The star beings to collapse inward because the fewer particles cannot maintain the pressure needed to support the star�s outer layers. Without the necessary pressure, the star�s core contracts slightly under the weight of its outer layers. This collapse increases the pressure and temperature of the core causing the luminosity of the star's surface to increase. This increase in pressure on the layers just outside the core raises their temperature to the necessary point in which the outer layers of hydrogen begin fusion. This occurrence of hydrogen fusion in the outer layers of a collapsing star is called shell hydrogen burning. Shell hydrogen burning allows the star to remain in the main sequence for another few million years. Despite this struggle to remain on the main sequence, the supply of hydrogen for fusion into helium in a star�s inner most layers depletes, which takes the star from the main sequence and to the next step of solar fuel consumption. The next stage of solar fuel consumption starts when hydrogen burning in the core ceases and ignites hydrogen burning in the star's outer layers. When hydrogen burning ceases in the star�s core, it begins to collapse again. At this point, the star converts gravitational energy into thermal energy because it must maintain thermal equilibrium. Stars sustain thermal equilibrium within their interiors through the ignition of helium burning. The collapse of the outer hydrogen burning shell upon the core raises the temperature and pressure in the core and begins helium burning. Because the temperature and energy needed to ignite helium fusion is greater than that of hydrogen, the energy released by helium fusion in the star�s core is greater than needed to support the weight of the outer layer. This excess energy expands the star�s outer layers beyond its previous radius and star�s volume increases. A star going through this stage of fuel consumption (collapse and expansion) is a Red Giant. The following diagram shows the dramatic expansion of a main sequence star as it begins helium fusion. Our Sun today compared to its future Red Giant self 1 AU is approximately 150 million kilometers or the distance from the sun to the Earth's orbit The Life of a Red Giant : A Star in Old Age. In the final states of hydrogen fusion the hydrogen burning shell adds to the mass of the star�s helium core. This added mass and pressure increases the star�s temperature. Temperature within the core of stars greater than three solar masses soon exceeds the 100 million Kelvin degrees barrier, at which it reaches the temperature needed for the onset of helium fusion. Core helium burning reestablishes the thermal equilibrium needed to support the star�s outer layers preventing further gravitational contraction. With gravitational contraction halted, the heat and energy from the core helium burning again begins the expansion of the star's outer layers. The method by which a star begins core helium burning depends upon the mass of the star. In stars below three solar masses (low-mass stars), helium fusion begins in a more spectacular manner. Helium burning begins explosively and abruptly. This process of quick ignition of helium fusion is a helium flash. The universe does not limit helium flash to low-mass stars; Helium is present in red giants and red supergiants as well. Red supergiants and red giants share the same physical characteristics and properties, with size and luminosity the only major difference between them. To understand this size difference we use our sun as a reference point. Red giants have solar radii up to 10 to 100 times larger than our sun, while supergiants boast a solar radii 1000 times larger than that of our sun. A Low-Mass Stars� last gasp. Stellar evolution can end in several ways. After a long life, an aging star completes its red supergiant stage of evolution and shell helium burning begins. Since this shell of helium burning is thin, the star becomes unstable. This instability in the aging star increases the temperature and energy within the star, which thickens the shell of burning helium surrounding the helium-depleted core. The shell of burning helium thickens until it can support the pressures of the star�s outer layers. As the shell helium burning increases the temperature and energy of the surrounding outer layers of the star, it ignites the outer hydrogen shell into fusion through thermonuclear reactions. This process of outer hydrogen shell fusion is a thermal pulse. During thermal pulses, the star�s luminosity increases by a factor of ten. This final expansion and ignition of outer layers is the star�s final moments before death overcomes the star. In its final moments, a star ejects its outer layers emitting ultraviolet radiation. This emission of ultraviolet radiation ionizes the ejected gases, giving them a glow known as a planetary nebula. A planetary nebula, the death of a low mass star The ejection of stellar remnants is the low-mass star�s supernova. On Earth, we measure the effects of this supernova in the increased luminosity. After a low-mass star�s death (supernova), it often leaves behind material that forms new stellar bodies. This is the end of stars with low masses (less than four solar masses) because they cannot reach the appropriate t temperature levels to begin the fusion of carbon (C) into oxygen (O). These stars burn off or eject their remaining outer layers leaving only the stellar core behind. Stellar remnants, a star�s afterlife: white dwarfs, neutron stars and pulsars. Due to the enormous mass of this remaining stellar core, the core begins to collapse and fuse. As the stellar core collapses, it electrons become degenerate (a phenomenon, due to the quantum mechanical effects, whereby the pressure exerted by a gas does not depend on its temperature) and stop the gravitation collapse. This contracted stellar core is a white dwarf. White dwarf stars are approximately the size of our planet, but their mass and density is much greater. White dwarfs are so dense, when compared to the Earth, that a teaspoon of matter from a white dwarf would weigh 5.5 metric tons on the Earth �s surface! The universe places limits on the life of a white dwarf, familiar to the main sequence star from which it originates. The white dwarf glows for billions of year from the energy released from cooling thermal radiation. Eventually all the radioactive matter of a white dwarf cools until it reaches the temperature of surrounding space which is a few degrees above absolute zero. A Neutron Star Neutron Stars and Pulsars : Sometimes the core remnant of a low-mass star is too dense to form a white dwarf. In these cases, the core�s stellar collapse forms a neutron star. Neutron stars are incredibly dense spheres of degenerate neutrons (a gas in which all the allowed states for particles, electrons or neutrons, have been filled, thereby causing the gas to behave differently from ordinary gases). Some neutron stars have massive magnetic fields that sweep out beams of radiation into space, much like light beams from a lighthouse. Such �pulsating� neutron stars are called pulsars. This measurable �pulse� of radiation is where this type of neutron star derives its name. Astronomical models suggest that the properties of superconductivity and superfluity dominate the cores of neutrons stars. These models also suggest there is an upper l limit in the mass of neutron stars. These upper limits are in the range of the mass of the largest possible white dwarfs. The Little Bang : The death of high-mass stars. This death of a low mass star into a white dwarf contrast with the final stages of high-mass, main sequence stars (i.e., greater than four solar masses). Unlike low mass stars, high mass stars can extend their lives through the fusion of elements heavier than carbon. After the fusion of helium ends, high mass stars begin burning carbon as their next fuel consumption stage. Carbon burning begins when the star's core reaches temperatures greater than 600 million degrees Kelvin. The greater the mass of the original main sequence star the longer the star continues to burn heavier elements. Incredibly massive stars continue fusion from helium until they create iron (Fe) in their core. In this fusion process, these massive stars create neon (Ne), magnesium (Mg), oxygen (O), sulfur (S), silicon (Si) and finally iron (Fe). Each stage of fuel consumption parallels a raise in the star�s core temperature. Neon fusion begins when the star�s core temperature reaches one billion degrees Kelvin followed by oxygen at 1.5 billion degrees Kelvin. Silicon fusion does not begin until the star�s core temperature reaches an amazing 3 billion degrees Kelvin. As the star begins each new fuel consumption stage, its lifetime in each stage becomes shorter. The following table illustrates the amount of time an aging star spends in each fuel consumption stage. A H-R diagram (The red line represents the Main Sequence stars and their evolution toward the top right hand corner of the diagram). These massive stars continue to exist by burning the heavier nuclides (magnesium, silicon, phosphorous et al) until they reach the heaviest nuclide that fusion can produce, iron. The iron rich cores of these massive stars can reach sizes equivalent to the Earth. The size of the surrounding shells burning the various lighter elements (H, He, C, O) dwarf this iron core, as they reach sizes that would extend to the orbit of Jupiter. When the core of a high-mass star consists completely of iron, fusion can no longer take place. At this point, the mass of the daughter (one iron nuclide) is heavier than the total mass of the parent nuclides used to produce iron. With this difference in mass, the process begins to absorb energy instead of releasing it back into the star, whichunbalances thethermal equilibrium of the star. With fusion ceased and the thermal equilibrium unbalanced, a star�s only source of energy is from the contraction of the star�s outer layers. This contraction raises the star�s core temperature to 5 billion Kelvin and increases the pressure to the point that the iron core collapses upon itself. The incredible pressures that the collapse of the star�s iron core force many of the existing (but not all) protons and electrons to combine into neutrons. For a moment the star releases this abundance of neutrons as neutrinos. In this brief escape, many neutrinos bombard the iron core and combine with iron nuclides to form the elements heavier than iron. These escaping neutrinos and the electromagnetic forces that repel protons and neutrons force the star into a final expansion. A Star�s Eulogy, a Supernova. This final expansion sends a shock wave through the star�s outer layers until it reaches the star�s surface creating a supernova. This shock wave ejects all the material of the star, including the core into the cosmos leaving behind only a nebula of cooling gasses. A section of space before the supernova The same section of space after the supernova Stars too big for their britches : Black Holes. Sometimes high-mass stars are too massive to become white dwarfs or neutron stars. A high-mass star this massive also has the gravitational forces to prevent the escape of stellar matter through a supernova. Stars with this great of mass become black holes at the end of their stellar evolution. Black holes are the result of the overpowering weight of the star�s massive outer layers pressing inward from all sides on the core causing the rapid contraction of the dying star. With this great mass pressing inward, the star�s core can no longer support the star�s outer layer structure. The matter within a high-mass star in this process is so condensed that the gravity produced is strong enough to prevent the escape of light energy from the cooling radioactive material. Light cannot escape because the escape velocity at the star�s surface is greater than that of the speed of light. Gravity of this magnitude has profound effects upon the star�s shape and on the structure of the surrounding space. With the star collapsing in upon itself, the star changes from a spherical shape to a broader plane in space. The gravity from the center of this plane curves its surrounding space to create a whirlpool-like drain that absorbs the star itself and all surrounding matter. An Artist rendition of a black hole event horizon (notice the whirlpool like shape that leads to the blackhole's singularity). This plane at the mouth of the whirlpool-like drain is called an event horizon. The immense gravitational forces of the black hole compacts all the matter within the star after it moves past the event horizon of the newly formed black hole. The gravity of the black hole compresses the star�s matter, as well as any other matter captured by the black hole�s gravitational forces through its existence, to infinite density. This infinitely dense matter within the black hole is called a singularity. The gravitational pull of a singularity is so great that it pulls matter from all surrounding areas into the event horizon of the black hole. If light cannot escape black holes, how do we find them? Until recently black holes were only theoretical speculations, but with today�s technology we are now finding black holes by measuring x-rayemissions in space. As the black holes immense gravity captures surrounding matter, this matter accelerates toward the event horizon and singularity and releases high amounts of x- rays. These x-rays leave a trail through space as the black hole absorbs the matter into its singularity. X-ray trails leave �halos� around the black hole called accretion disks. Another method of detecting black holes in space is their effects on binary star systems. When a black hole occurs close to two stars, its gravity impacts their spatial relationships. Each star�s gravity has a calculable effect upon the other�s orbit, however the invisible black hole affects these orbits. An Artist rendition of a blackhole syphoning plasma gas from a nearby star (this acceleration of matter emits x-rays that form a accretion disk and make the blackhole detectable [x-ray emissions represented by the arrows]). To contrast the idea of black holes being places in the universe with infinite lifetimes where no matter (including light) ever escapes, the universe holds a few exceptions. The British physicist Stephen W. Hawking has proven that black holes do in fact have definite lifetimes. As black holes reach the end of their life, they begin to evaporate. In the final stages of this evaporation, the black hole reverses itself and pours matter back out into the universe. When a black hole begins to eject its matter, it is a white hole. With this transformation in the life of a black hole, the universe appears to maintain the fundamental universal energy-matter law with this process. While the black hole, to physics' laws, is an unbalancing factor in universal laws, the white hole exists to restore this matter and balance. Our Two-cents--Please, no change Stellar evolution is the necessary fundamental building block and distributive method of most common elements in the universe. Within the interior of stars, fusion creates new elements from the basic elements (H, He). While this process takes billions of years as measured by human standards, the life of a star is minor in comparison to the age of the universe. Much like its life span, the outcome of one star�s life is insignificant. When we add the efforts of one star, however, to the production of the billions of stars that have existed, exist and will exist again they are invaluable to the creation of life. Without the elements produced by stellar evolution (the processes within the interiors of stars), Carbon-based life as we know it would not be possible. The next time you are star gazing, remember, you are a product of the stars. HERTZSPRUNG-RUSSELL DIAGRAM The Hertzsprung-Russell diagram (HR diagram) is one of the most important tools in the study of stellar evolution. Developed independently in the early 1900s by Ejnar Hertzsprung and Henry Norris Russell, it plots the temperature of stars against their luminosity (the theoretical HR diagram), or the colour of stars (or spectral type) against their absolute magnitude (the observational HR diagram, also known as a colour- magnitude diagram). Depending on its initial mass, every star goes through specific evolutionary stages dictated by its internal structure and how it produces energy. Each of these stages corresponds to a change in the temperature and luminosity of the star, which can be seen to move to different regions on the HR diagram as it evolves. This reveals the true power of the HR diagram – astronomers can know a star’s internal structure and evolutionary stage simply by determining its position in the diagram. The Hertzsprung-Russell diagram the various stages of stellar evolution. By far the most prominent feature is the main sequence (grey), which runs from the upper left (hot, luminous stars) to the bottom right (cool, faint stars) of the diagram. The giant branch and supergiant stars lie above the main sequence, and white dwarfsare found below it. Credit: R. Hollow, CSIRO. This Hertzsprung-Russell diagram shows a group of stars in various stages of their evolution. By far the most prominent feature is the main sequence, which runs from the upper left (hot, luminous stars) to the bottom right (cool, faint stars) of the diagram. The giant branch is also well populated and there are many white dwarfs. Also plotted are the Morgan-Keenan luminosity classes that distinguish between stars of the same temperature but different luminosity. --> There are 3 main regions (or evolutionary stages) of the HR diagram: 1. The main sequence stretching from the upper left (hot, luminous stars) to the bottom right (cool, faint stars) dominates the HR diagram. It is here that stars spend about 90% of their lives burning hydrogen into helium in their cores. Main sequence stars have a Morgan-Keenan luminosity class labelled V. 2. red giant and supergiant stars (luminosity classes I through III) occupy the region above the main sequence. They have low surface temperatures and high luminosities which, according to the Stefan-Boltzmann law, means they also have large radii. Stars enter this evolutionary stage once they have exhausted the hydrogen fuel in their cores and have started to burn helium and other heavier elements. 3. white dwarf stars (luminosity class D) are the final evolutionary stage of low to intermediate mass stars, and are found in the bottom left of the HR diagram. These stars are very hot but have low luminosities due to their small size. The Sun is found on the main sequence with a luminosity of 1 and a temperature of around 5,400 Kelvin. Astronomers generally use the HR diagram to either summarise the evolution of stars, or to investigate the properties of a collection of stars. In particular, by plotting a HR diagram for either a globular or open cluster of stars, astronomers can estimate the age of the cluster from where stars appear to turnoff the main sequence (see the entry on main sequence for how this works). ……………………………………………………………………………………. Luminosity From Wikipedia, the free encyclopedia For other uses, see Luminosity (disambiguation). Image of galaxy NGC 4945 showing the huge luminosity of the central few star clusters, suggesting there is an AGNlocated in the center of the galaxy. In astronomy, luminosity is the total amount of energy emitted by a star, galaxy, or other astronomical object per unit time. It is related to the brightness, which is the luminosity of an object in a given spectral region.  In SI units luminosity is measured in joules per second or watts. Values for luminosity are often given in the terms of the luminosity of the Sun, L⊙. Luminosity can also be given in terms of magnitude: the absolute bolometric magnitude (Mbol) of an object is a logarithmic measure of its total energy emission rate. Contents [hide] 1Measuring luminosity 2Stellar luminosity 3Radio luminosity 4Magnitude 5Luminosity formulae 6Magnitude formulae 7See also 8References 9Further reading 10External links Measuring luminosity Hertzsprung–Russell diagramidentifying stellar luminosity as a function of temperature for many stars in our solar neighborhood. In astronomy, luminosity is the amount of electromagnetic energy a body radiates per unit of time.  When not qualified, the term "luminosity" means bolometric luminosity, which is measured either in the SI units, watts, or in terms of solar luminosities (L☉). A bolometeris the instrument used to measure radiant energy over a wide band by absorption and measurement of heating. A star also radiates neutrinos, which carry off some energy (about 2% in the case of our Sun), contributing to the star's total luminosity. The IAU has defined a nominal solar luminosity of 3.828×1026 W to promote publication of consistent and comparable values in units of the solar luminosity.  While bolometers do exist, they cannot be used to measure even the apparent brightness of a star because they are insufficiently sensitive across the electromagnetic spectrum and because most wavelengths do not reach the surface of the Earth. In practice bolometric magnitudes are measured by taking measurements at certain wavelengths and constructing a model of the total spectrum that is most likely to match those measurements. In some cases, the process of estimation is extreme, with luminosities being calculated when less than 1% of the energy output is observed, for example with a hot Wolf-Rayet star observed only in the infra-red. Bolometric luminosities can also be calculated using a bolometric correction to a luminosity in a particular passband.  The term luminosity is also used in relation to particular passbands such as a visual luminosity of K- band luminosity. These are not generally luminosities in the strict sense of an absolute measure of radiated power, but absolute magnitudes defined for a given filter in a photometric system. Several different photometric systems exist. Some such as the UBV or Johnson system are defined against photometric standard stars, while others such as the AB system are defined in terms of a spectral flux density. Stellar luminosity A star's luminosity can be determined from two stellar characteristics: size and effective temperature.  The former is typically represented in terms of solar radii, R⊙, while the latter is represented in kelvins, but in most cases neither can be measured directly. To determine a star's radius, two other metrics are needed: the star's angular diameter and its distance from Earth, often calculated using parallax. Both can be measured with great accuracy in certain cases, with cool supergiants often having large angular diameters, and some cool evolved stars having masers in their atmospheres that can be used to measure the parallax using VLBI. However, for most stars the angular diameter or parallax, or both, are far below our ability to measure with any certainty. Since the effective temperature is merely a number that represents the temperature of a black body that would reproduce the luminosity, it obviously cannot be measured directly, but it can be estimated from the spectrum. An alternative way to measure stellar luminosity is to measure the star's apparent brightness and distance. A third component needed to derive the luminosity is the degree of interstellar extinction that is present, a condition that usually arises because of gas and dust present in the interstellar medium (ISM), the Earth's atmosphere, and circumstellar matter. Consequently, one of astronomy's central challenges in determining a star's luminosity is to derive accurate measurements for each of these components, without which an accurate luminosity figure remains elusive. Extinction can only be measured directly if the actual and observed luminosities are both known, but it can be estimated from the observed colour of a star, using models of the expected level of reddening from the interstellar medium. In the current system of stellar classification, stars are grouped according to temperature, with the massive, very young and energetic Class O stars boasting temperatures in excess of 30,000 K while the less massive, typically older Class M stars exhibit temperatures less than 3,500 K. Because luminosity is proportional to temperature to the fourth power, the large variation in stellar temperatures produces an even vaster variation in stellar luminosity.  Because the luminosity depends on a high power of the stellar mass, high mass luminous stars have much shorter lifetimes. The most luminous stars are always young stars, no more than a few million years for the most extreme. In the Hertzsprung–Russell diagram, the x-axis represents temperature or spectral type while the y-axis represents luminosity or magnitude. The vast majority of stars are found along the main sequence with blue Class 0 stars found at the top left of the chart while red Class M stars fall to the bottom right. Certain stars like Deneb and Betelgeuse are found above and to the right of the main sequence, more luminous or cooler than their equivalents on the main sequence. Increased luminosity at the same temperature, or alternatively cooler temperature at the same luminosity, indicates that these stars are larger than those on the main sequence and they are called giants or supergiants. Blue and white supergiants are high luminosity stars somewhat cooler than the most luminous main sequence stars. A star like Deneb, for example, has a luminosity around 200,000 L ⊙, a spectral type of A2, and an effective temperature around 8,500 K, meaning it has a radius around 203 R⊙. For comparison, the red supergiant Betelgeuse has a luminosity around 100,000 L ⊙, a spectral type of M2, and a temperature around 3,500 K, meaning its radius is about 1,000 R⊙. Red supergiants are the largest type of star, but the most luminous are much smaller and hotter, with temperatures up to 50,000 K and more and luminosities of several million L⊙, meaning their radii are just a few tens of R⊙. An example is R136a1, over 50,000 K and shining at over 8,000,000 L⊙ (mostly in the UV), it is only 35 R⊙. …………………………………….. Chandrasekhar limit The Chandrasekhar limit is the maximum mass of a stable white dwarf star. The currently accepted value of the Chandrasekhar limit is about 1.4 M☉(2.765×1030 kg). White dwarfs resist gravitational collapse primarily through electron degeneracy pressure (compare main sequence stars, which resist collapse through thermal pressure). The Chandrasekhar limit is the mass above which electron degeneracy pressure in the star's core is insufficient to balance the star's own gravitational self-attraction. Consequently, a white dwarf with a mass greater than the limit is subject to further gravitational collapse, evolving into a different type of stellar remnant, such as a neutron star or black hole. Those with masses under the limit remain stable as white dwarfs. Collapse is not inevitable: most white dwarfs explode rather than undergo collapse. The limit was first indicated in papers published by Wilhelm Anderson and E. C. Stoner in 1929. It was named after Subrahmanyan Chandrasekhar, the Indian astrophysicist who independently discovered and improved upon the accuracy of the calculation in 1930, at the age of 20, in India. The limit was initially ignored by the community of scientists because such a limit would logically require the existence of black holes, which were considered a scientific impossibility at the time.