Electromagnetic Radiation - Assignment Example

Written Assignment 2 Please answer each of the standard essay questions in this assignment in a minimum of one typed page The Seeds textbook s: "To an astronomer, nothing is so precious as starlight." (p. 72) Today we recognize that this starlight is electromagnetic radiation. List the following: A. Each part of this radiation used by an astronomer. The starlight coming from deep space is useful to an astronomer because it is the only information that we dispose about the physical phenomena of the Universe. The electromagnetism is the mathematical tool of an astronomer. The electromagnetism says that starlight is an electromagnetic wave. An electromagnetic wave is the couples in space-time between an electric field and magnetic field. This physical perturbation of and electric field and a magnetic field in the Universe is well described by the Maxwell’s equations. The electric field and magnetic field in an electromagnetic ware are orthogonal each other, as well as to the direction of propagation. Both electric and magnetic fields in an electromagnetic wave are variable in time. The Maxwell’s equation says that any electromagnetic wave has a velocity propagation called c. Nowadays, we know that the speed of light in vacuum has a value near to 300,000 km/s. The speed of light in any other medium is less than in the vacuum. Physics tell us that the light has energy flux given by the Poything vector, then is has momentum. Furthermore, the light has radiation pressure. The plane wave model of the Maxwell’s equation is a trivial solution, however there exist other models such as spherical waves. An electromagnetic wave is described also by its wavelength; this is related to its period T and its speed, c. B. What instrument (telescope, etc.) works best with each type of radiation. The human eye is sensible to the Sun’s light this is due to man evolution. The human have invented instruments able to detect almost any electromagnetic radiation coming from deep Universe, in fact is coming from anywhere. Additionally, the man has designed observational methods able to increase the quality of surveys. The optical telescope works best in the visible part of the electromagnetic spectrum. The visible window is around between 7000 A and 4 000 A. The radio astronomer uses huge radio dish antennas, these instruments work best at large electromagnetic wavelength, i.e. > 109 A. When the radiation is between 109 and 106 A the astronomers use bollometers in order to detect the millimeter/submillimetre wavelength. At infrared level between 106A and 7000 A they use CCD detectors at very low temperature. In order to detect ultraviolet waves i.e. between 4000A and 10A the astronomers use solid state devices. The X-Rays between 10A - 0.1A the detection is given by collision between particles. The most energetic particles are called Gamma rays it has < 0.1 A, thus the detection is using external electromagnetic fields i.e. Lorentz force. C. What celestial objects we study and what we learn about them from the radiation they emit. It is well known that when a charge is moving this produce electromagnetic radiation or when a molecule or atom is excited it passes to its lowest energy state, it produce electromagnetic radiation. D. What special instruments work in conjunction with telescopes to advance our studies of "starlight"? The telescope is the ideal instrument to study the starlight; however in order to record such observations the astronomers use the CCD (Charge Coupled Device) detector. The detector consists of a surface made up of light sensitive silicon diodes, arranged in a rectangular array of image elements or pixels. The largest cameras can have as many as 4096x4096 pixels. Another useful instrument to study starlight is called diffraction grid. With this, we can register stellar spectrum. 2. The ultimate key to our understanding the universe is our knowledge of the atom. A. Illustrate with an example the difference between an atom and an ion. An atom is described as a nucleus surrounded by a swarm of electrons. The nucleus consist of Z protons, each having a charge e+ and N electrically neutral neutrons; Z is the charge number of the atom and A = Z + N is its mass number. A neutral atom has as many electrons with charge e- as protons. On the hand, when an atom gains or losses an electron (or electrons) resulting in a charged atom. For example: Name Atom Ion Description Hydrogen H H+ Hydrogen gains an electron, becomes negatively charged. Sodium Na Na+ Sodium gain an electron, becomes negatively charged. Chlorine Cl Cl- Chlorine losses an electron, becomes positively charge. B. Describe two ways an atom can be excited. An atom can become excited in two ways, by absorbing some energy from a source of electromagnetic radiation of by colliding with some other particle, i.e. another atom. C. Why should photons emitted by a hotter material have an average shorter wavelength? If we assume that a hot material has a black body radiation, the maximum intensity obeys the Wien’s law. λmax T = 2.8978 m-3 K. As we can see, increasing the temperature of the material, we have shorter wavelength. D. Atoms produce spectra. Distinguish between a continuous, a bright-line, and an absorption spectrum by describing how each is formed. Continuous emission spectra can originate in recombination and free-free transitions. In recombination, an atom captures a free electron whose energy is not quantized; in free-free transitions, both initial and final states are unquantized. Thus the emission line can have any frequency whatsoever. Similarly, ionizations and free-free transitions can give rise to a continuous absorption spectrum. When the pressure of hot gas is increased, the spectral lines begin to broaden. At high pressure, atoms bump into each other more frequently, and the close neighbors disturb the energy levels. When the pressure is high enough, the lines begin to overlap. Thus the spectrum of hot gas at high pressure is continuous. Electric fields also broaden spectral lines. In liquids and solids, the atoms are more densely packed than in gaseous substances. Their mutual perturbations broaden the energy levels, producing continuous spectrum. 3. How was the spectral classification system arrived at? Relate its construction to the Balmer series and explain how a stars spectral class can give us clues to its temperature, motion, and chemical composition. The following research has been taken from Kartunnen 1996. Physical properties of stars can be derived directly from studies of their spectra. There are two main important spectral classifications: The Harvard Spectral Classification and The Yerkes Spectral Classification. The Harvard classification is based on lines that are mainly sensitive to stellar temperature, rather than to gravity or luminosity. Important lines are the hydrogen Balmer lines, the lines of neutral helium, the iron lines, the H and K doublet of ionized calcium at 396.8 and 393.3 nm, the G band due to the CH molecule and some metals around 431 nm, the neutral calcium line at 422.7 nm and the lines of titanium oxide (TiO). The main types in the Harvard classification are denoted by capital letters. They were initially ordered in alphabetical order, but subsequently it was noticed that they could be ordered according to temperature. With the temperature decreasing towards the right the sequence is: C O – B – A – F – G – K – M S Additional notations are Q for novae, P for planetary nebulae and W for Wolf-Rayet stars. The class C consists of the earlier types R and N. The spectral classes C and S represent parallel branches to types G-M, differing in their surface chemical composition. The spectral classes are divided into sublclases denoted by the numbers 0…9; sometimes decimals are used i.e. B0.5. The main characteristics of the different classes are: Type O: Blue stars, surface temperature 20,000-35,000 K. Spectrum with lines from multiply ionized atoms, e.g. He II, C III, N III, O III, Si IV, He I visible, H I lines weak. Type B: Blue-white stars, surface temperature about 15,000 K. He II lines have disappeared, He I (403 nm) lines are strongest at B2. Then get weaker and have disappeared at type B9, The K line of Ca II becomes visible at type B3. H I lines getting stronger. O II, Si II and Mg II lines visible. Type: A. White stars, surface temperature about 9,000 K. The H I lines are very strong at A0 and dominate the whole spectrum, then get weaker. H and K lines of Ca III getting stronger. He I no longer visible. Neutral metal lines begin to appear. Type F: Yellow-white stars, surface temperature about 7,000 K. H I lines getting weaker, H and K of Ca II getting stronger. Many other metal lines, e.g. Fe I, Fe II, Cr II, Ti II, clear and getting stronger. Type G: Yellow stars like the Sun, surface temperature about 5,500 K. The H I lines still getting weaker, H and K lines very strong, strongest at G0. Metal lines getting stronger. G band clearly visible. CN lines seen in giant stars. Type K: Orange-yellow stars, surface temperature about 4,000 K. Spectrum dominated by metal lines. H I insignificant. Ca I 422.7 nm clearly visible. Strong H and K lines and G band. TiO bands becomes visible at K5. Type M: Red stars, surface temperature about 3,000 K. TiO band getting stronger. Ca I 422.7 nm very strong. Many neutral metal lines. Type C: Carbon stars, previously R and N. Very red stars, surface temperature about 3,000 K. Strong molecular bands, e.g. C2, CN and CH. No TiO bands. Line spectrum like the K and M Types. Type S: Red low-temperature stars (about 3,000 K). Very clear ZrO bands. Also other molecular bands, e.g. YO, LaO and TiO. In order to understand how the strengths of the spectral lines are determined by temperature: The neutral helium lines at 402.6 nm and 447.2 nm, these lines are due to absorption by excited atoms, and that at high temperature is required to produce any appreciable excitation. As the stellar temperature increases, more atoms are in the required excited state, and the strength of the helium lines increase. When the temperature becomes even higher, helium begins to be ionized, and the strength of the neutral helium lines begins to decreases. In a similar way one can understand the variation with temperature of other important lines, such as the hydrogen Balmer lines. The Yerkes classification is based on the visual scrutiny of slit spectra with a dispersion of 11.5 nm/mm. It is carefully defined on the basis of standard stars and the specification of luminosity criteria. Six luminosity classes are distinguished. Ia most luminous supergiant Ib less luminous supergiant II luminous giants III normal giants IV subgiants V main sequence stars. The luminosity class is determined form spectral lines that depend strongly on the stellar surface gravity, which is closely related to the luminosity. The gravitational acceleration is calculated as: g = G M / R2. In consequence, the gas density and pressure in the atmosphere of a giant is much smaller than in a dwarf star. For spectral type B – F, the lines of neutral hydrogen are deeper and narrower for stars of higher luminosities. The reason for this is that metal ions give rise to a fluctuating electrical field near the hydrogen atoms. This filed leads to shift in the hydrogen energy levels, appearing as a broadening of the lines. The effect becomes stronger as the density increases. Thus the hydrogen lines are narrow in absolutely bright stars, and become broader in main sequence stars and even more so in white dwarfs. The lines from ionized elements are relatively stronger in high-luminosity stars. This is because a higher density makes it easier for electrons and ions to recombine to neutral atoms. On the other hand, the rate of ionization is essential determined by the radiation field, and is not appreciably affected by the gas density. Thus a given radiation field can maintain a higher degree of ionization in stars with more extended atmospheres. For example, in the spectral classes F – G, the relative strengths of the ionized strontium (Sr II) and the neutral iron (Fe I) lines can be used as a luminosity indicator. Both lines depend on the temperature in roughly the same way, but the Sr II lines become relatively much stronger as the luminosity increases. Giant stars are redder than dwarfs of the same spectral type. The spectral type is determined from the strengths of spectral lines, including ion lines. Since these are stronger in giants, a giant will be cooler, and thus also redder, than a dwarf of the same spectral type. There is a strong cyanogens (CN) absorption band in the spectral of giant stars, which is almost totally absent in dwarfs. This is partly a temperature effect, since the cooler atmospheres of giants are more suitable for the formation of cyanogens. 4. Describe the sun in terms of its structure an activity. Include in your discussion the following terms: sunspots, prominences, solar flares, and solar wind. The following research has been taken from Kartunnen 1996. The Sun is the main source of energy of the solar system. The Sun is the nearest star to us. The Sun is a typical main sequence star. The Sun is a yellow star G2 V type with surface temperature around 5,785 K. Its principal properties are: Mass m = 1.989 x 1030kg Radius R = 6.960 x 108 m Mean density = 1409 kgm-3 Central density = 1.6 x 105 kg m-3 Luminosity L = 3.9 x 1026 W Effective temperature Te = 5785 K Central temperature, Tc = 1.5 x 107 K Absolute bolometric magnitude, Mbol = 4.72 Absolute visual magnitude, Mv = 4.79 Spectral class, G2 V Color indices, B-V = 0.62, U-B = 0.10 Surface chemical composition, X = 0.71 (Hydrogen), Y(Helium) = 0.27, Z(other elements) = 0.02 Rotational period, at the equator = 25 days, at 60o latitude = 29 days The most realistic solar model give us additional information: The energy is produced by the pp chain in a small central region. 99% of the solar energy is produced within a quarter of the solar radius. The Sun produces energy of rate of 4 x 1026 W. The Atmosphere: The solar atmosphere is divided into the photosphere and the chromospheres. Outside the actual atmosphere, the corona extends much further outwards. The photosphere: The innermost layer of the atmosphere is the photosphere, which is only 300-500 km thick. The photosphere is the visible surface of the Sun, where the density rapidly increases inwards, hiding the interior from sight. The temperature at the inner boundary of the photosphere is 8,000 K and at the outer boundary 4,500 K. Near the edge of the solar disc, the lines of sight enters the photosphere at a very small angle and never penetrates to large depths. Near the edge one therefore only sees light form the cooler, higher layers. For this reason, the edges appear darker; this phenomenon is known as limb darkening. Both the continuum spectrum and the absorption lines are formed in the photosphere, but the light in the absorption lines comes from higher layers and therefore the lines appears dark. The solar convection is visible on the surface as the granulation, an uneven, constantly changing granular pattern. At the bright centre of each granule, gas is rising upward, and at the darker granule boundaries, it is sinking down again. The size of a granule seen from the earth is typically 1’’, corresponding to about 1,000 km on the solar surface. There is also a larger scale convection called supergranulation in the photosphere. The cells of the supergranulation may be about 1´ diameters. The observed velocities in the supergranulation are mainly directed along the solar surface. The Chormosphere: Outside the photosphere there is a layer, perhaps about 500 km thick, where the temperature increases from 4,500 K to about 6,000 K, the chromospheres. Outside this layer, there is a transition region of a few thousand kilometers, where the chromospheres gradually goes over into the corona. In the outer parts of the transitions region, the kinetic temperature is already about 106 K. Normally the chromospheres is not visible, because its radiation is so much weaker than that of the photosphere. However, during total solar eclipses, the chromospheres shines into view for a few seconds at both ends of the total phase, when the Moon hides the photosphere completely. The chromospheres than appears as a thin reddish sickle or ring. During eclipses, the chromospheric spectrum, called the flash spectrum, can be observed. It is an emission line spectrum with more than 3,000 indentified lines. Brightest among these are the lines of hydrogen, helium and certain metals. The Corona: The corona is also best seen during total solar eclipses. It then appears as a halo of light extending out to a few solar radii. The surface brightness of the corona is about that of the full moon, and it is therefore difficult to see next to the bright photosphere. The inner part of the corona, the K corona, has a continuous spectrum formed by the scattering of the photospheric light by electrons. Further out, a few solar radii from the surface is the F corona, which has a spectrum showing Fraunhofer absorption lines. The light of the F corona is sunlight scattered by dust. The helium was discovered in the corona before it was known on Earth. This is because was detected helium highly ionized. In fact, in the corona there is thirteen times ionized iron. This is due by entire corona has to have a temperature of about a millions degrees. In spite of its high temperature, the coronal gas is so diffuse that the total energy stored in it is small. It is constantly streaming outwards, gradually becoming a solar wind, which carries a flux of particles away from the Sun. The gas lost in this way is replaced with new material from the chromospheres. Near the Earth, the density of the solar wind is typically 5-10 particles/cm3 and its velocity about 500 km/s. The mass loss of the Sun due to the solar wind is about 10-13 solar masses per year. Solar activity: The clearest visible sign of solar activity are the sunspot. The existence of sunspot has been known for long. When Galilei started to use the telescope he reported small spot on the Sun. A sunspot looks like a ragged hole in the solar surface. In the interior of the spot is a dark umbra and around it, a less dark penumbra. The surface temperature in a sunspot is about 1,500 K below that of its surrounding, which explains the dark color of the sports. The diameter of a typical sunspot is about 10,000 km and its lifetime is from a few days to several months, depending on its size. The larger spots are more likely to be long-lived. Sunspot often occurs in pairs or in larger groups. By following the motions of the spots, the period of rotation of the Sun can be determined. The variations on the number of sunspot have been followed for almost 250 years. The frequency of spots is described by Zurich sunspot number Z: Z = C(S+10G) where S is the number of spot and G the number of spot groups visible at a particular time. C is a constant depending on the observer and the conditions of observation. The figure 1 shows the Zurich sunspot number from 1700 to 1992. It is clear that spots vary with an average period of about 11 years. Figure 1: It shows the Zurich sunspot number from 1700 to 1992. The number of sunspot and spot groups varies with period of about 11 years. The periodic variation in the number of sunspot reflects a variation in the general solar magnetic field. At the beginning of new activity cycles spots first begin to appear about 40o latitude. As the cycles advances, the spots move closer to the equator. The characteristic pattern in which spots appear, shown in figure 2 is known as the butterfly diagram. Spots of the next cycle begin to appear while those of the old one are still present near the equator. Spots belonging to the new cycle have a polarity opposite that of the old ones. Sin the field is thus reversed between consecutive 11 years cycles the complete period of solar magnetic activity is 22 years. Figure 2: At the beginning of an activity cycles, sunspot appear at high latitudes. As the cycles advances the spot move towards the equator. Other activity: The Sun shows several other types of surface activity: faculae and plages; prominences; flares. The faculae and plages are local bright regions in the photosphere and chromospheres, respectively. Observations of the plages are made in the hydrogen Hα or the calcium K lines. The plages usually occur where new sunspots are forming, and disappear when the spots disappear. Apparently they are caused by the enhanced heating of the chromospheres in strong magnetic fields. The prominences are among the most spectacular solar phenomena. They are glowing gas masses in the corona, easily observed near the edge of the Sun. There are several types of prominences: The quiescent prominences, where the gas is slowly sinking along the magnetic field lines; loop prominences, connected with magnetic field loops in sunspots; and the rarer eruptive prominences, where gas is violently thrown outwards. The temperature of prominences is about 10,000 – 20,000 K . In Hα photographs of the chromospheres, the prominences appears as dark filaments against the solar surface. The flare outburst are among the most violent forms of solar activity. They appear as bright flashes, lasting from one second to just under an hour. In the flares, a large amount of energy stored in the magnetic field is suddenly released. The detailed mechanism is not yet known. Flares can be observed at all wavelengths. The hard x-ray emission of the Sun may increase a hundred times during a flare. Several different types of flares are observed at radio wavelengths. The emission of solar cosmic ray particles also rises. 5. If Earth and the humans who inhabit it can be affected by solar activity, describe how. Since the Sun is the main source of energy in the solar system, particularly on Earth. The solar flares give rise to disturbances on the Earth by affecting electrical systems, especially satellites. The x-ray cause changes in the ionosphere, which affect short-wave radio communications. The flare particles give rise to strong aurora when they enter the Earth’s magnetic field a few days after the outburst. There still debate about connection between solar activity and impact on human life on Earth. However, on May 5th 2010 a solar eruption came from spotless region, producing a bright coronal mass ejection and came to an end on May 9th. The SOHO satellite captured the solar eruption, see figure X (see SOHO, 2010). Figure X. It shows solar eruption on May 5th, by NASA. In short the solar activity affects the Earth’s magnetic field, but there still poor connection between solar activity and Earth events like earthquakes, volcanoes, hurricanes, floods and storms. The solar activity may affect technology on the following way: Ionosphere Variations: Wireless signal reflection, propagation, attenuation. Communication satellite signal interference, scintillation Interference with geographical prospecting Source of electrical currents in the Earth. Power distribution systems Long communication cables, land and ocean Pipelines. Radiation Solar cell damage Semiconductor device damage and failure Misoperation of semiconductor devices. Spacecraft charging, surface and interior materials Astronaut safety Airline passenger safety Magnetic field variations Attitude control of spacecraft Compasses Solar radio burst Excess noise levels in wireless communications systems Atmosphere Low altitude satellite drag Attenuation and scatter of wireless signals. Solar activity impact on Earth climate: There have been suggestions that climate is related o the appearance and disappearance of sunspot. This effect and other solar activity offer the best explanation for the observed rise in average global temperature over last century (Panagiotis, 2010). References Panagiotis K. M. 2010 “The stormy Sun affecting the human life and the techology” Vol 32. http://www.the-eggs.org. SOHO 2010, http://sohowww.nascom.nasa.gov, accessed on 9 Nov. 2010 Karttunen H., Kroger P., Oja A.h, Poutanen M. and Donner K. J., 1996 “Fundamental Astronomy” Springer-Verlag Berlin Heidelberg New York. Read More