how atomic spectra're produced n how they're used in astronomy to determine the speed and temperature of stars
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Well if you look at the spectrum of a star you'll see a smooth continuum as well as the spectral lines.  Those lines are either absorbtion lines or emission lines of such things as Hydrogen, Helium, Calcium, Iron and a whole heap of others and we know the wavelength of those lines very well so by comparing what we measure them to be with what we have determined in the lab we can calculate how fast the star is moving towards or away us (radial velocity).\n\nThen once we know the radial velocity we just measure where the peak in the continuum is, correct for the radial velocity to find what the wavelength where the emission actually peaks and then since the star is very close to a blackbody we can calculate the temperature (a blackbody of a given temperature will peak at a given wavelength).
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Isolated atoms can absorb and emit packets of electromagnetic radiation having discrete energies dictated by the detailed atomic structure of the atoms. When the corresponding light is passed through a prism or spectrograph it is separated spatially according to wavelength.\n\nThe corresponding spectrum may exhibit a continuum, or may have superposed on the continuum bright lines (an emission spectrum) or dark lines (an absorption spectrum). \n\nThus, emission spectra are produced by thin gases in which the atoms do not experience many collisions (because of the low density). The emission lines correspond to photons of discrete energies that are emitted when excited atomic states in the gas make transitions back to lower-lying levels.\n\nAn absorption spectrum occurs when light passes through a cold, dilute gas and atoms in the gas absorb at characteristic frequencies; since the re-emitted light is unlikely to be emitted in the same direction as the absorbed photon, this gives rise to dark lines (absence of light) in the spectrum.
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Astronomers determined the size and shape of our galaxy and measured distances to other stars using the “parallax method” (and eventually other methods using the star’s light spectrum).  Unfortunately, these methods do little to help us find out the exact distance to our very own star, the Sun.  Surprisingly, to get the distance to the sun, astronomers first had to determine the exact distance to Venus!\n\nWith the invention of radar, scientists were able to measure the distance to Venus very precisely.  By timing how long it took the radar beam to travel at the speed of light to Venus and back, the total distance to the object can be determined from the equation:\n\nDistance = (Speed of Light) X (Total Time)/2\n\nThe reason the total time is divided by two is to get just the distance from the Earth to the object.\nEarth-Sun Distance = (Earth-Venus Distance) / (Cos q)\n\nUsing trigonometry, astronomers now know that the “Astronomical Unit” = 92,955,820.5 miles (149,597,892 kilometers).  This incredible degree of accuracy is possible because the speed of light is known very precisely and extremely accurate clocks are used to time the radar reflection.  Radar can’t be used to determine the distance to the Sun directly because the Sun doesn’t have a solid surface to reflect the radar efficiently.
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When performing star counts, astronomers consider many different categories which have created been to classify a few stars that have been well studied. One of the hopes of studying the results of star counts is to discover new categories. Different counts typically seek to categorize stars for only a few of the qualities listed below, and determine how common each considered quality is and how stars of that kind are distributed.\n\nTemperature. In astronomy, temperature is usually shown using the letter codes O B A F G K M L T running from blue (type O) through white (type F) to red (type M). Types L and T are used for brown dwarfs, whose 'colors' are in the infrared. \nSize. Size is usually designated by Roman numerals I (supergiants) through V (dwarfs). \nAge. Stars are usually grouped into Population I (young) and Population II (old). \nLocation. In the Milky Way Galaxy the groups are described as thin disk, thick disk, central bulge, and halo. \nMultiplicity. Most stars appear to be members of double star or triple star, or even double-double star systems. Our own sun appears to be rare for not having a companion star. \nThere are many finer subdivisions in all of the above categories.\nIn which category you want me to count the Stars.
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Stars are balls of gas, They burn by converting lighter elements into heavier elements.  The lightest element is hydrogen, but stars make heavier elements as the burn, so really old stars have elements as dense a Iron!\n\nHeres a link to the periodic table:\n\nhttp://www.dayah.com/periodic/\n\nAll the elements between hydrogen and iron could be in a star.\n\nBut the "Easy" answer is to say Hydrogen and Helium!
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I am not quite sure what your asking.\n\nWhen the star is forming, it starts out as a gigantic cloud of mostly Hydrogen. As the stellar material coalesces, the material begins to heat up and the pressure at the center of what will become the new star increases dramatically.  The material becomes denser and is pulled together by its own gravity, heating up as the pressure rises.  Gravity is working to pull the star together, and eventually the ball of gas becomes so massive, that the pressure and heat at its core are sufficient to undergo nuclear fusion.  In order to counteract the force of gravity, the Hydrogen at the stars core undergoes fusion, in the process, producing Helium and a lot of energy.  It is this energy produced from nuclear fusion at the stars core that keeps the star from collapsing in on itself.
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It isn't measured, it is calculated. The speed of light in miles per second is multiplied by the number of seconds in a year to get the length of a light year in miles. \n\nMeasuring the speed of light can be done in many ways, and is not as hard to do as it was before electronics were invented. Early methods relied on flashes of light from rotating mirrors miles apart. Today laser pulses are timed very accurately with electronics over short distances. Also, amazingly, Maxwell's equations calculate the speed of light theoretically from a number of other unrelated electric and magnetic measurements.\n\nMeasuring the distance to remote stars is another topic, but the results are usually expressed in light years because the numbers are just inconveniently large when expressed in miles.
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In the 76th edition it is Chapter 10 of the CRC - "Line Spectra of the Elements".  You may need to convert the numbers given to the units you desire.
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Because of its proximity to the earth, and because it is such a typical star, the sun is a unique resource for the study of stellar phenomena. No other star can be studied in such detail. The star closest to the sun is 4.3 light-years (4 × 1013 km/2.5 × 1013 mi) away. To observe features on its surface comparable to those that can be seen routinely on the sun would require a telescope almost 30 km (18.6 mi) in diameter. Such a telescope, moreover, would have to be put into space to avoid distortions caused by the earth's atmosphere.\n\n    The Sun is an ordinary G2 star, one of more than 100 billion stars in our galaxy.\n\n \n\n        diameter:    1,390,000 km.\n        mass:        1.989e30 kg\n        temperature: 5800 K (surface)\n                     15,600,000 K (core)\n   The Sun is by far the largest object in the solar system. It contains more than 99.8% of the total mass of the Solar System (Jupiter contains most of the rest). \n   The Sun is personified in many mythologies: the Greeks called it Helios and the Romans called it Sol. \n\n   The Sun is, at present, about 75% hydrogen and 25% helium by mass (92.1% hydrogen and 7.8% helium by number of atoms); everything else ("metals") amounts to only 0.1%. This changes slowly over time as the Sun converts hydrogen to helium in its core. \n\n \n\n   The outer layers of the Sun exhibit differential rotation: at the equator the surface rotates once every 25.4 days; near the poles it's as much as 36 days. This odd behaviour is due to the fact that the Sun is not a solid body like the Earth. Similar effects are seen in the gas planets. The differential rotation extends considerably down into the interior of the Sun but core of the Sun rotates as a solid body. \n\n   Conditions at the Sun's core (approximately the inner 25% of its radius) are extreme. The temperature is 15.6 million Kelvin and the pressure is 250 billion atmospheres. At the centre of the core the Sun's density is more than 150 times that of water. \n\n   The Sun's energy output (3.86e33 ergs/second or 386 billion billion megawatts) is produced by nuclear fusion reactions. Each second about 700,000,000 tons of hydrogen are converted to about 695,000,000 tons of helium and 5,000,000 tons (=3.86e33 ergs) of energy in the form of gamma rays. As it travels out toward the surface, the energy is continuously absorbed and re-emitted at lower and lower temperatures so that by the time it reaches the surface, it is primarily visible light. For the last 20% of the way to the surface the energy is carried more by convection than by radiation. \n\n    The surface of the Sun, called the photosphere, is at a temperature of about 5800 K. Sunspots are "cool" regions, only 3800 K (they look dark only by comparison with the surrounding regions). Sunspots can be very large, as much as 50,000 km in diameter. Sunspots are caused by complicated and not very well understood interactions with the Sun's magnetic field. \n\n   A small region known as the chromosphere lies above the photosphere. \n\n    The highly rarefied region above the chromosphere, called the corona, extends millions of kilometers into space but is visible only during eclipses (left). Temperatures in the corona are over 1,000,000 K. \n\n   The Sun's magnetic field is very strong (by terrestrial standards) and very complicated. Its magnetosphere (also known as the heliosphere) extends well beyond Pluto. \n\n    In addition to heat and light, the Sun also emits a low density stream of charged particles (mostly electrons and protons) known as the solar wind which propagates throughout the solar system at about 450 km/sec. The solar wind and the much higher energy particles ejected by solar flares can have dramatic effects on the Earth ranging from power line surges to radio interference to the beautiful aurora borealis.  \n\n   The solar wind has large effects on the tails of comets and even has measurable effects on t
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The distances to stars closer than about 150 light years can be measured directly by use of parallax.  The position of the star against background stars is measured at times 1/2 year apart, and then it is a simple mathematical procedure to determine the distance (you have a side and two angles of a triangle--the angular distance the star appears to move at it's distance over 1/2 year equals 186 million miles-the diameter of the orbit of the earth).  \n\nStars farther than that cannot be measured directly.  But, by use of knowledge about how bright a star should be (on average) for it's spectral type (which loosely is governed by its mass) an estimate of its intrinsic brightness can be made, and then by looking at how bright it appears, it's distance can be estimated.\n\nCepheid variables are used, as well, but mostly to determine the distances to groups of stars and of other galaxies.  Incidentally, the North Star (Polaris, or Alpha Ursa Minor) is a cepheid variable.\n\nNovae and supernovae tend to have average intrinsic brightnesses, so their apparent brightness in, say, a distant galaxy, tells us (within limits) it's distance.\n\nRed shift is used for objects outside our galaxy because the amount of red shift within our galaxy would be extremely small (it increases with distance) if it were not totally overcome by the gravity of the galaxy.
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milky way is our own galaxy.sThe compositions of stars are determined through spectroscopy. Spectroscopy is the study of something using spectra. Recall from the electromagnetic radiation chapter that a spectrum is what results when you spread starlight out into its individual colors. By noting what absorption lines (or sometimes, emission lines) are present and their strengths, you can find out a tremendous amount of information. Stars have absorption lines patterns similar to the Sun. This means that they are composed mostly of hydrogen and helium with traces of other elements. \n\nFrom these absorption lines you learn some important things beside the star's composition: \n\nStructure of stars: From the simple fact that you see absorption lines in most stellar spectra, you know that the stars must have a hot dense part that produces a continuous spectrum and an outer layer, or atmosphere, made of cooler, low density gas. The general trend is density and temperature of stars decreases as the distance from the star's center increases. The hot dense part is also gaseous because of the extreme temperatures. Stars have no molten rock in them like the interiors of some of the planets. \nmass of stars:\nTo determine the masses of stars, Kepler's third law is applied to the motions of binary stars---two stars orbiting a common point. The greater the combined mass of the two stars, the greater the gravity acceleration is, and, therefore, the smaller their orbital period. A majority of the several hundred billion stars in the Galaxy are in a system with two or more stars orbiting each other. Usually the binary stars are spectroscopic binary stars. A spectroscopic binary system is two stars orbiting a common point at too great a distance away from us to resolve the two stars individually, but whose binary nature is indicated in the periodic shift of their spectral lines as they orbit around each other. Spectroscopic binary stars are used because (a) there are a lot more far away stars than nearby ones and (b) more importantly, you can easily measure their speeds from the doppler shifted lines. \ncolour of star:\nHot stars appear bluer than cooler stars. Cooler stars are redder than hotter stars. The ``B-V color index'' is a way of quantifying this using two different filters; one a blue (B) filter that only lets a narrow range of colors or wavelengths through centered on the blue colors, and a ``visual'' (V) filter that only lets the wavelengths close to the green-yellow band through. \nA hot star has a B-V color index close to 0 or negative, while a cool star has a B-V color index close to 2.0. Other stars are somewhere in between. Here are the steps to determine the B-V color index: \n\nMeasure the apparent brightness (flux) with two different filters (B, V). \nThe flux of energy passing through the filter tells you the magnitude (brightness) at the wavelength of the filter. \nCompute the magnitude difference of the two filters, B - V. \n\nok thats all,
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First, our sun has a long time to go before it could possibly become a supernova, not for billions and billions of years, so don't worry about it, if you were worrying. OK?\n\nNow, the way a star burns is by fusion, which is building up heavier elements by fusing atoms of lighter elements together. The typical start starts out as hydrogen, and it fuses the hydrogens together into helium. \n\nWhat happens when a star runs out of hydrogen? It starts fusing the heliums together into heavier elements. In order to begin this process, the star has to contract, to squeeze everything together to make the core temperature hotter. \n\nNotice that when the star began to run out of hydrogen, it cooled off a little. So the temperature of the star fluctuates. Its size fluctuates too, since the heat energy partly counteracts the gravity that holds it together.\n\nSo, the star cools off a little, it contracts, and we call this a pulsating or variable star. (This is somewhat of an oversimplification; there are different types of variable stars. But I am ignoring them for purposes of this illustration.) Then the star begins fusing helium, and heats up again. The star expands. \n\nThe outer surface of the star explodes away in what we call a nova-- not a supernova, yet. Some of the gas rushes away and may form what is called a nebula, a cloud. some of thee nebulas give off pretty colors, because, just like a neon sign, the gas in them is excited by the energy of the star and glows.\n\nSo, after repeated cycles of this, and bazillions of years, the star is exhausted. When the star gets iron in its core, there is no turning back. You can't fuse iron into anything. \n\nBut, you say, what about those elements that are heavier than iron, like gold and mercury and uranium? Ah, I'm so glad you asked about that. That is where the supernova comes in. \n\nA supernova is the death of a star. A nova flares up and puts out lots of energy, but the supernova puts out as much energy as the star might have put out over its entire lifetime! The last time there was a supernova in our galaxy (the Crab Nebula was formed) in 1054 AD. The Chinese were the only ones whe were doing much astronomical observation then. \n\nBut it was visible in broad daylight! From LIGHT YEARS away!\n\nWhat happens is this: the star begins to contract one final time. Most of the star blows away, making a really, really impressive nebula. All the elements that were built up in the final explosion-- all the heavy elements beyond iron-- get flung out into space, where they might meet with other elements and form planets and things.\n\nWhat happens to the rest of the star? One of three things, depending on how big it is. If it is little, it may form a white dwarf, a little star that is not so big but is now mostly hydrogen and helium and will last a long time. I think they are supposed to expand into red giants when they are about to burn out, but I don't remember for sure. \n\nIf the star is medium sized, it may become a neutron star. It is no longer buring with fusion like a regular star. The nuclear particles in its core are squeezed together so that protons and electrons have lost their identities and everything is neutrons. This is as solid and dense as matter can get, neutronium. It has an incredibly strong magnetic field, bazillions of times anything on earth.\n\nThe star's axis of rotation may not line up with the magnetic poles (just the way Earth's magnetic poles are not right at the true north and true south poles.) If this hapens, and the magnetic poles are pointed in the right direction, then every time it points our way, we see a little 'flash' of magnetic energy, and we call these 'pulsars'. The first time anybody detected them, they thought they were radio signals from people in space, because they were so regularly-spaced. It is a good thing there are no pulsars near Earth, because the flashes of intense magnetic fields would fry us lik
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try these sites: http://www.phy.ornl.gov/astrophysics/data/cf88/\n\nhttp://www.llnl.gov/science_on_lasers/07AstSpace/AS-F_Ther.html\n\nhttp://arxiv.org/abs/astro-ph/9801028\n\nI hope this helps....\n\nDon
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A star is a massive body of plasma in outer space that is currently producing or has produced energy through nuclear fusion. The most familiar and closest star to the Earth is the Sun. Unlike a planet, from which most light is reflected, a star emits light because of its intense heat. Scientifically, stars are defined as self-gravitating spheres of plasma in hydrostatic equilibrium, which generate their own energy through the process of nuclear fusion. Stellar astronomy is the study of stars.
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The stars themselves don't.  It is just an illusion.  The miles of air in our atmosphere between our eyes and the star changes in temperature and humidity which makes the star look like it is flickering when it really glows at a constant rate.
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Spectroscopy is the branch of science (Physics) concerned with production and analysis of spectra (analogous range of radiations in order of wavelength), produced by a spectrophotometer.\n\nIn emission spectroscopy, a sample is normally excited by thermal energy. Atoms in the sample absorb thermal energy causing outer orbital electrons to become excited from their ground state to higher orbital energy level. After a short lifetime (10{+-}{+6} to 10{+-}{+9}s), the excited electrons return to the ground state. Simultaneously, electromagnetic radiations (normally in the form of light in the UV-visible region) are emitted. The emitted radiations are analysed by means of spectrograph, which separates various wavelengths.\n\nIn Plasma Emission Spectroscopy, the exciting medium/ agent is plasma. The term `Plasma' in Physics means a very hot/ high temperature ionised gas, sometimes described as the fourth state of matter (the other three being solid, liquid and gas). In this method of sample analysis, a `Plasma Source' is the basic component. (Plasma Emission Spectroscopy is similar to Optical Emission Spectroscopy in principles but differs in mechanism. In OES, the sample excitation is done by electric arc or spark mechanism in the temperature range of 3000-5000 degrees Celcius. In PES, the temperature range is 7000-9000 {+o}C.\n\nThe sample solution or the analyte solution is supplied to the `Plasma Source' by a `Nebulizer' assembly. On reaching the high temperature (7000-9000 degrres Celsius) plasma, the solution attains this high temperature and all types of molecular bonds break down. Free atoms or ions (electrically charged atoms) are produced, which emit their characteristic spectra. A `Monochromator' assembly is in position after `Plasma Source' for isolation of the desired analyte line (or spectrum). The intensity of an analyte line is directly proportional to the concentration of that particular analyte (element being analysed) in the sample solution. The `detector & read-out' assembly gives us the concentration of the particular element in the sample. Modern instruments being computer controlled, the analysis is very fast and needs little operator skill. (Depending on the technique of production, plasma may be of three classes — viz., D.C. argon plasma, Microwave plasma and Inductively coupled plasma (ICP). The ICP is the most widely used due to mechanical simplicity and operational advantages.\n\nTherefore, plasma Emission Spectroscopy is the method of analysis of various elements in a sample, by passing the sample/ analyte solution through a plasma (high temperature ionised gas) source, thereby exciting the outer orbital electrons of the analyte with simultaneous emission of electromagnetic radiations (in the form of lights), which are analysed by means of a spectrograph, which separates various wavelengths of analytes in the sample.
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STARS AND NEBULAE.\n\nMost of the things you can see in the night sky are stars: a few thousand are visible to the unaided eye. A star is a hot ball of mostly hydrogen gas; the Sun is an example of a typical, ordinary star. Gravity keeps the gas from evaporating into space, and pressure due to the star's high temperature and density keeps the ball from shrinking. In the core of the star, the temperature and densities are high enough to sustain nuclear fusion reactions, and the energy produced by these reactions works its way to the surface and radiates into space as heat and light. When the fuel for the fusion reactions is depleted, the structure of the star changes. The process of building up heavier elements from lighter ones by nuclear reactions, and adjusting the internal structure to balance gravity and pressure, is called stellar evolution. \nLooking at a star through a telescope can tell us many of its important properties. The color of a star tells us its temperature, and the temperature depends on some combination of the star's mass and evolutionary phase. Usually, observations also allow us to find the luminosity of the star, or the rate at which it radiates energy as heat and light. \nAll stars visible to the unaided eye are in our galaxy, the Milky Way. The Milky Way is system of roughly 100 billion stars, along with a great deal of interstellar material. The galaxy is shaped like a flattened disk embedded in a round, faint halo. Gravity keeps the stars from drifting away, and the motions of the stars keep the system from falling in on itself. The Milky Way has no distinct edge - the distribution of stars tapers off gradually with increasing distance from the center. The SDSS detects stars more than a million times fainter than can be seen with the unaided eye, far enough to see the outer structure of the Milky Way. \nStar Colors and Luminosities: The H-R Diagram\nAstronomers often make plots of the luminosities of stars compared to their colors. The first such charts were made in the early twentieth century, when astronomers took spectra of thousands of stars, then placed the spectra into a sequence based on the appearance of several different features. The classes of stars were given different letters as names; in order of decreasing temperature, they are OBAFGKMLT. The hottest stars are class O, while the coolest are class T. Two astronomers, Ejnar Hertzsprung and Henry Norris Russell, independently started to wonder what would happen if they compared stars' luminosities to their spectral classifications (or, alternatively, to their temperatures). They knew that some stars were hotter and more luminous than the sun, and that others were cooler and less luminous. Hertzsprung and Russell found that 90% of the stars fell into a narrow band they termed the "main sequence." Today, we call this type of plot a Hertzsprung-Russell (or H-R) diagram. \nAstronomers further divide each letter into ten categories from 0 to 9, with 0 being the hottest star in that spectral class and 9 being the coolest. So, a B1 star is hotter than a B2 star, and a B9 star is hotter than an A0 star. Our own sun is at the "hot" end of the G range with a surface temperature of 5,770° K; therefore, it's called a G2 star. \nSpectral classes give us a convenient way to refer to most types of stars - from O0, the hottest at over 25,000° K, down to T9, the coolest failed stars at 1,000° K. Each spectral type of star is associated with a specific color: all G-type stars like the sun are yellow, all M-type stars appear red, and all B-type stars are bluish. Thus, this system of letters and numbers tells us where a star is on the horizontal (color or temperature) axis of the H-R diagram. \nTwo alternative systems tell us approximately where on the vertical (luminosity) axis a star can be found. One system uses absolute magnitude, a measure of luminosity derived from the apparent magnitude, or brightness of the star