Spectral analysis of stars. Spectral analysis in astronomy. Significance for cosmology
PAGE_BREAK - The total radiation of the Sun is determined by the illumination it creates on the Earth's surface - about 100 thousand lux when the Sun is at its zenith. Outside the atmosphere at the average distance of the Earth from the Sun, the illumination is 127 thousand lux. The intensity of the sun's light is 2.84 10527 candles. The amount of energy that comes in one minute to an area of \u200b\u200b1 cm, placed perpendicular to the sun's rays outside the atmosphere at the average distance of the Earth from the Sun, is called the solar constant. The total radiation power of the Sun is 3.83 10526 watts, of which about 2 10 517 watts fall on the Earth, the average brightness of the Sun's surface (when observed outside the Earth's atmosphere) is 1.98 1059 nt, the brightness of the center of the Sun's disk is 2.48 1059 nt ... The brightness of the solar disk decreases from the center to the edge, and this decrease depends on the wavelength, so that the brightness at the edge of the solar disk for light with a wavelength of 3600 A is 0.2 brightness of its center, and for 5000 A - about 0.3 brightness of the center disk of the sun. At the very edge of the Sun's disk, the brightness falls 100 times in less than one arc second, so the edge of the Sun's disk looks very sharp.
The spectral composition of the light emitted by the Sun, that is, the distribution of energy in the center of the Sun (after taking into account the influence of absorption in the earth's atmosphere and the influence of Fraunhofer lines), in general terms corresponds to the distribution of energy in the radiation of a black body with a temperature of about 6000 K. However, in some parts of the spectrum there are noticeable deviations. The maximum energy in the solar spectrum corresponds to a wavelength of 4600 A. The solar spectrum is a continuous spectrum, which is superimposed on more than 20 thousand absorption lines (Fraunhofer lines). More than 60% of them have been identified with spectral lines of known chemical elements by comparing the wavelengths and relative intensities of the absorption line in the solar spectrum with laboratory spectra. The study of Fraunhofer lines provides information not only about the chemical composition of the solar atmosphere, but also about the physical conditions in those layers in which these or those absorptions are formed. The predominant element in the sun is hydrogen. The number of helium atoms is 4–5 times less than that of hydrogen. The number of atoms of all other elements taken together is at least 1000 times less than the number of hydrogen atoms. The most abundant of them are oxygen, carbon, nitrogen, magnesium, iron and others. In the solar spectrum, one can also identify lines belonging to certain molecules and free radicals: OH, NH, CH, CO and others.
Magnetic fields on the Sun are measured mainly by the Zeeman splitting of absorption lines in the solar spectrum. There are several types of magnetic fields on the Sun. The total magnetic field of the Sun is small and reaches a strength of 1 of this or that polarity and changes over time. This field is closely related to the interplanetary magnetic field and its sectorial structure.
Magnetic fields associated with solar activity can reach intensities of several thousand in sunspots. The structure of magnetic fields in active regions is very confused; magnetic poles of different polarities alternate. There are also local magnetic regions with field strengths of hundreds outside sunspots. Magnetic fields penetrate both the chromosphere and the solar corona.
Magnetogasdynamic and plasma processes play an important role on the Sun.
At a temperature of 5000-10000 K, the gas is sufficiently ionized, its conductivity is high, and due to the huge scale of solar phenomena, the importance of electromechanical and magnetomechanical interactions is very great.
Sun atmosphere
The Sun's atmosphere is formed by the outer layers accessible to observation. Almost all of the Sun's radiation comes from the lower part of its atmosphere, called the photosphere. Based on the equations of radiative energy transfer, radiant and local thermodynamic equilibrium, and the observed radiation flux, it is possible to theoretically construct a model of the temperature and density distribution with depth in the photosphere. The photosphere is about three hundred kilometers thick, with an average density of 3 104–5 kg / m. The temperature in the photosphere decreases with the transition to more outer layers, its average value is about 6000 K, at the border of the photosphere is about 4200 K. The pressure varies from 2 1054 to 1052 n / m.
The existence of convection in the sub-photospheric zone of the Sun is manifested in the uneven brightness of the photosphere, its visible granularity - the so-called granulation structure. The granules are bright spots of a more or less round shape. The size of the granules is 150 - 1000 km, the life time is 5 - 10 minutes, individual granules can be observed within 20 minutes. Sometimes granules form clusters up to 30 thousand kilometers in size. The granules are 20-30% brighter than the intergranular gaps, which corresponds to an average temperature difference of 300 K. Unlike other formations, granulation on the Sun's surface is the same at all heliographic latitudes and does not depend on solar activity. The velocities of chaotic motions (turbulent velocities) in the photosphere are, according to various definitions, 1–3 km / sec. Quasiperiodic oscillatory motions in the radial direction were found in the photosphere. They occur on sites measuring 2–3 thousand kilometers with a period of about five minutes and a velocity amplitude of about 500 m / s. After several periods, the fluctuations in this place die out, then they can arise again. Observations have also shown the existence of cells in which movement occurs in the horizontal direction from the center of the cell to its boundaries. The speed of such movements is about 500 m / sec. The dimensions of the cells - supergranules - are 30-40 thousand kilometers. The position of the supergranules coincides with the cells of the chromospheric grid. At the boundaries of supergranules, the magnetic field is enhanced.
Supergranules are thought to reflect, at a depth of several thousand kilometers, below the surface of convective cells of the same size. Initially, it was assumed that the photosphere gives only continuous radiation, and absorption lines are formed in the reversing layer located above it. Later it was found that both spectral lines and a continuous spectrum are formed in the photosphere. However, to simplify mathematical calculations when calculating spectral lines, the concept of a reversing layer is sometimes used.
Sun spots and torches are often observed in the photosphere.
Sun spots
Sunspots are dark formations consisting, as a rule, of a darker core (shadow) and a partial shade surrounding it. The diameters of the spots reach two hundred thousand kilometers. Sometimes the spot is surrounded by a light border.
Very scarlet spots are called pores. The lifetime of stains is from several hours to several months. The spectrum of sunspots contains even more absorption lines and bands than the spectrum of the photosphere; it resembles the spectrum of a star of the spectral type KO. The displacement of lines in the spectrum of spots due to the Doppler effect indicates the movement of matter in spots - outflow at lower levels and inflow at higher levels, the speed of movement reaches 3 thousand m / sec. From comparisons of the line intensities and the continuous spectrum of sunspots and the photosphere, it follows that the spots are colder than the photosphere by 1–2 thousand degrees (4500 K and below). As a result, against the background of the photosphere, the spots appear dark, the brightness of the core is 0.2 - 0.5 brightness of the photosphere, the brightness of penumbra is about 80% of the photospheric. All sunspots have a strong magnetic field, reaching an intensity of 5 thousand oesters for large sunspots. Typically, spots form groups that, in terms of their magnetic field, can be unipolar, bipolar and multipolar, that is, they contain many spots of different polarity, often united by a common penumbra. Groups of sunspots are always surrounded by torches and floccules, prominences, solar flares sometimes occur near them, and formations in the form of rays of helmets and fans are observed in the solar corona above them - all this together forms an active region on the Sun. The average annual number of observed sunspots and active regions, as well as the average area occupied by them, changes with a period of about 11 years.
This is an average value, while the duration of individual cycles of solar activity ranges from 7.5 to 16 years. The largest number of spots simultaneously visible on the surface of the Sun changes more than twice for different cycles. Most of the sunspots are found in the so-called royal zones, stretching from 5 to 30 ° heliographic latitude on both sides of the solar equator. At the beginning of the solar activity cycle, the latitude of the location of the sunspots is higher, and at the end of the cycle, it is lower, and at higher latitudes, sunspots of a new cycle appear. Bipolar sunspot groups are more often observed, consisting of two large sunspots - the head and the next, having the opposite magnetic polarity, and several smaller ones. The head sunspots have the same polarity throughout the entire solar cycle; these polarities are opposite in the northern and southern hemispheres of the Sun. Apparently, the spots are depressions in the photosphere, and the density of matter in them is less than the density of matter in the photosphere at the same level.
Torches
In active regions of the Sun, torches are observed - bright photospheric formations visible in white light mainly near the edge of the Sun's disk. Torches usually appear before the spots and exist for some time after they disappear. The area of \u200b\u200bflare areas is several times larger than the area of \u200b\u200bthe corresponding group of spots. The number of torches on the solar disk depends on the phase of the solar cycle. The maximum contrast (18%) of the torches is near the edge of the solar disk, but not at the very edge. In the center of the Sun's disk, the torches are practically invisible, their contrast is very small. Torches have a complex fibrous structure, their contrast depends on the wavelength at which observations are made. The temperature of the torches is several hundred degrees higher than the temperature of the photosphere, the total radiation from one square centimeter exceeds the photospheric one by 3 - 5%. The torches seem to rise somewhat above the photosphere. The average duration of their existence is 15 days, but can reach almost three months.
Chromosphere
Above the photosphere is a layer of the Sun's atmosphere called the chromosphere. Without special telescopes, the chromosphere is visible only during total solar eclipses as a pink ring surrounding the dark disk in those minutes when the Moon completely covers the photosphere. Then the spectrum of the chromosphere can be observed. At the edge of the Sun's disk, the chromosphere appears to the observer as an uneven strip, from which separate denticles - chromospheric spicules - protrude. The diameter of the spicules is 200-2000 kilometers, the height is about
It is believed that spicules are formed at the boundaries of the cells of the chromospheric grid. The density in the chromosphere decreases with increasing distance from the center of the Sun. The number of atoms in one cube. centimeter varies from 10515 0 near the photosphere to 1059 in the upper part of the chromosphere. The study of the spectra of the chromosphere led to the conclusion that in the layer where the transition from the photosphere to the chromosphere occurs, the temperature passes through a minimum and, as the height above the base of the chromosphere increases, it becomes equal to 8-10 thousand Kelvin, and at an altitude of several thousand kilometers it reaches 15-20 thousand Kelvin.
It was found that in the chromosphere there is a chaotic movement of gas masses with speeds up to 15 1053 m / s. In the chromosphere, torches in active regions are visible as light formations, usually called floccules. In the red line of the hydrogen spectrum, dark formations called fibers are clearly visible. At the edge of the Sun's disk, filaments protrude beyond the disk and are observed against the sky as bright prominences. Most often, filaments and prominences are found in four zones located symmetrically relative to the solar equator: polar zones north of + 40 ° and south of -40 ° heliographic latitude and low latitude zones around √ (30 °) at the beginning of the solar activity cycle and √ (17 °) at the end cycle. Fibers and prominences of low-latitude zones show a well-pronounced 11-year cycle, their maximum coincides with the maximum of sunspots.
In high-latitude prominences, the dependence on the phases of the solar activity cycle is less pronounced; the maximum occurs two years after the maximum of sunspots.
The fibers, which are calm prominences, can reach the length of the solar radius and exist for several revolutions of the Sun. The average height of prominences above the surface of the Sun is 30-50 thousand kilometers, the average length is 200 thousand kilometers, and the width is 5 thousand kilometers. According to the research of A.B. North, all prominences by the nature of the movement can be divided into 3 groups: electromagnetic, in which movements occur along ordered curved trajectories - the lines of force of the magnetic field; chaotic, in which disordered turbulent movements prevail (speeds of the order of 10 km / s); eruptive, in which the substance of the initial calm prominence with chaotic movements is suddenly ejected with an increasing speed (reaching 700 km / sec) away from the Sun. The temperature in the prominences (fibers) is 5-10 thousand Kelvin, the density is close to the average density of the chromosphere. Fibers, which are active, rapidly changing prominences, usually change dramatically in a few hours or even minutes. The form and character of the movements in prominences are closely related to the magnetic field in the chromosphere and solar corona.
The solar corona is the outermost and most rarefied part of the solar atmosphere, extending over several (more than 10) solar radii. Until 1931, the corona could be observed only during total solar eclipses in the form of a silvery-pearl glow around the Sun's disk covered by the Moon. The details of its structure stand out well in the crown: helmets, fans, coronal rays and polar brushes. After the invention of the coronagraph, the solar corona was also observed outside eclipses. The general shape of the corona changes with the phase of the solar activity cycle: in the years of the minimum, the corona is strongly elongated along the equator, and in the years of the maximum it is almost spherical. In white light, the surface brightness of the solar corona is a million times less than the brightness of the center of the solar disk. Its glow is formed mainly as a result of the scattering of photospheric radiation by free electrons. Almost all atoms in the corona are ionized. The concentration of ions and free electrons at the base of the corona is 1059 particles per cm. The heating of the corona is carried out similarly to the heating of the chromosphere. The greatest energy release occurs in the lower part of the crown, but due to its high thermal conductivity, the crown is almost isothermal - the temperature drops outward very slowly. The outflow of energy in the corona occurs in several ways.
In the lower part of the corona, the downward transfer of energy due to thermal conductivity plays the main role. The loss of energy is caused by the departure of the fastest particles from the corona. In the outer parts of the corona, most of the energy is carried away by the solar wind - a flow of coronal gas, the speed of which increases with distance from the Sun from several km / s at its surface to 450 km / s at the distance of the Earth. The temperature in the corona exceeds 1056 K. In the active layers of the corona, the temperature is higher, up to 1057 K. Above the active regions, so-called coronal condensations can form, in which the concentration of particles increases tens of times. Part of the radiation inside the corona is the radiation lines of multiply ionized atoms of iron, calcium, magnesium, carbon, oxygen, sulfur and other chemical elements. They are observed both in the visible part of the spectrum and in the ultraviolet region. In the solar corona, radio emission from the Sun in the meter range and X-rays are generated, which are amplified many times in active regions. Calculations have shown that the solar corona is not in equilibrium with the interplanetary medium.
From the corona to interplanetary space, particle streams propagate, forming the solar wind. Between the chromosphere and the corona, there is a relatively thin transition layer, in which there is a sharp rise in temperature to values \u200b\u200bcharacteristic of the corona. The conditions in it are determined by the flow of energy from the corona as a result of heat conduction. The transition layer is the source of most of the sun's ultraviolet radiation.
The chromosphere, transition layer, and corona give all the observed radio emission from the Sun. In active regions, the structure of the chromosphere, corona, and transition layer changes. This change, however, is still not well understood.
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--PAGE_BREAK - In active regions of the chromosphere, sudden and relatively short-term increases in brightness are observed, visible at once in many spectral lines. These bright formations last from several minutes to several hours. They are called solar flares (formerly called chromospheric flares). Flares are best seen in the hydrogen line light, but the brightest are sometimes seen in white light. The solar flare spectrum contains several hundred emission lines of various elements, neutral and ionized. The temperature of those layers of the solar atmosphere that give rise to luminescence in chromospheric lines (1–2) x1054 K, in higher layers - up to 1057 K. The density of particles in a flare reaches 10513 -10514 in one cubic centimeter. The area of \u200b\u200bsolar flares can reach 10515 m. Typically, solar flares occur near rapidly developing sunspot groups with a complex magnetic field. They are accompanied by the activation of fibers and floccules, as well as the release of matter. A flash produces a lot of energy (up to 10521 - 10525 joules).
It is assumed that the energy of a solar flare is initially stored in a magnetic field and then quickly released, which leads to local heating and acceleration of protons and electrons, causing further heating of the gas, its glow in different parts of the electromagnetic radiation spectrum, and the formation of a shock wave. Solar flares give a significant increase in the ultraviolet radiation of the Sun, are accompanied by bursts of X-ray radiation (sometimes very powerful), bursts of radio emission, and the ejection of high-energy carpuscles up to 10510 eV. Occasionally, bursts of X-ray radiation are observed without intensifying the glow in the chromosphere.
Some flares (they are called proton flares) are accompanied by especially strong fluxes of energetic particles - cosmic rays of solar origin.
Proton flares pose a danger to astronauts in flight, colliding with the atoms of the spacecraft shell, since energetic particles generate X-rays and gamma radiation, and sometimes in dangerous doses.
The level of solar activity (the number of active regions and sunspots, the number and intensity of solar flares, etc.) changes with a period of about 11 years. There are also weak fluctuations in the magnitude of the 11-year cycle peaks with a period of about 90 years. On Earth, an 11-year cycle can be traced in a number of phenomena of organic and inorganic nature (disturbances of the magnetic field, auroras, disturbances of the ionosphere, changes in the growth rate of trees with a period of about 11 years, established by the alternation of the thickness of annual rings, etc.). Terrestrial processes are also influenced by separate active regions on the Sun and the short-term, but sometimes very powerful flares occurring in them. The lifetime of a separate magnetic region on the Sun can be up to one year. The disturbances caused by this region in the magnetosphere and the upper atmosphere of the Earth are repeated 27 days later (with the period of rotation of the Sun observed from the Earth). The most powerful manifestations of solar activity - solar (chromospheric) flares occur irregularly (more often near periods of maximum activity), their duration is 5–40 minutes, rarely several hours. The energy of a chromospheric flare can reach 10525 joules, of the energy released during a flare, only 1–10% falls on electromagnetic radiation in the optical range. In comparison with the total radiation of the Sun in the optical range, the energy of the flare is not large, but the short-wave radiation of the flare and the electrons generated during the flares, and sometimes the solar cosmic rays, can make a noticeable contribution to the X-ray and carpuscular radiation of the Sun. During periods of increased solar activity, its X-ray radiation doubles in the range of 30-10 nm, in the range of 10-1 nm by 3-5 times, in the range of 1-0.2 nm more than a hundred times. As the radiation wavelength decreases, the contribution of active regions to the total radiation of the Sun increases, and in the last of the indicated ranges, practically all radiation is due to active regions. Hard X-rays with wavelengths less than 0.2 nm appear in the solar spectrum only for a short time after flares. In the ultraviolet range (wavelength 180–350 nm), the solar radiation for an 11-year cycle changes by only 1–10%, and in the range of 290–2400 nm it remains practically constant and amounts to 3.6 10526 watts.
The constancy of the energy received by the Earth from the Sun ensures the stationarity of the Earth's heat balance. Solar activity does not significantly affect the energy of the Earth as a planet, but individual components of the emission of chromospheric flares can have a significant impact on many physical, biophysical and biochemical processes on Earth.
Active regions are a powerful source of corpuscular radiation. Particles with energies of about 1 keV (mainly protons) propagating along the lines of force of the interplanetary magnetic field from active regions intensify the solar wind. These amplifications (gusts) of the solar wind are repeated after 27 days and are called recurrent. Similar flows, but of even greater energy and density, arise during flares. They cause the so-called sporadic disturbances of the solar wind and reach the Earth in time intervals from 8 hours to two days. High energy protons (from 100 MeV to 1 GeV) from very strong "proton" flares and electrons with energies of 10–500 keV, which are part of solar cosmic rays, come to the Earth tens of minutes after the flares; somewhat later, those of them come that fell into the "traps" of the interplanetary magnetic field and moved along with the solar wind. Short-wave radiation and solar cosmic rays (at high latitudes) ionize the earth's atmosphere, which leads to fluctuations in its transparency in the ultraviolet and infrared ranges, as well as to changes in the propagation conditions of short radio waves (in some cases, violations of short-wave radio communication are observed).
The intensification of the solar wind caused by the flare leads to the compression of the Earth's magnetosphere from the solar side, increased currents at its outer boundary, partial penetration of solar wind particles deep into the magnetosphere, replenishment of high-energy particles of the Earth's radiation belts, etc. These processes are accompanied by fluctuations in the intensity of the geomagnetic field (magnetic storm), auroras and other geophysical phenomena reflecting the general disturbance of the Earth's magnetic field. The influence of active processes on the Sun (solar storms) on geophysical phenomena is carried out both by short-wave radiation and through the Earth's magnetic field. Apparently, these factors are the main ones for physicochemical and
biological processes. It has not yet been possible to trace the entire chain of connections leading to the 11-year periodicity of many processes on Earth, but the accumulated vast factual material leaves no doubt about the existence of such connections. So, a correlation was established between the 11-year cycle of solar activity and earthquakes, crop yields, the number of cardiovascular diseases, etc. These data indicate the constant action of solar-terrestrial connections.
Observations of the Sun are carried out using small to medium-sized refractors and large mirror telescopes, in which most of the optics are stationary, and the sun's rays are directed into the horizontal or tower installation of the telescope using one or two moving mirrors. A special type of solar telescope - an extra-eclipse coronagraph - has been created. Inside the coronagraph, the sun is darkened with a special opaque screen. In the coronagraph, the amount of scattered light decreases many times, therefore, outside the eclipse, the outermost layers of the Sun's atmosphere can be observed. Solar telescopes are often equipped with narrow-band light filters that allow observations in the light of a single spectral line. Neutral light filters with variable transparency along the radius have also been created, which make it possible to observe the solar corona at a distance of several solar radii. Usually large solar telescopes are equipped with powerful spectrographs with photographic or photoelectric fixation of spectra. The spectrograph can also have a magnetograph - a device for studying the Zeeman splitting and polarization of spectral lines and determining the magnitude and direction of the magnetic field on the Sun. The need to eliminate the washing effect of the earth's atmosphere, as well as studies of solar radiation in the ultraviolet, infrared and some other regions of the spectrum, which are absorbed in the Earth's atmosphere, led to the creation of orbital observatories outside the atmosphere, which make it possible to obtain spectra of the Sun and individual formations on its surface outside the earth's atmosphere ...
The path of the sun among the stars
Every day, rising from the horizon in the eastern side of the sky, the Sun passes across the sky and hides again in the west. For residents of the Northern Hemisphere, this movement occurs from left to right, for southerners from right to left. At noon, the Sun reaches its greatest height, or, as astronomers say, climaxes. Noon is the upper culmination, and there is also the lower one - at midnight. In our mid-latitudes, the sun's lower culmination is not visible, as it occurs below the horizon. But beyond the Arctic Circle, where the Sun sometimes does not set in summer, one can observe both the upper and lower culminations.
At the geographic pole, the diurnal path of the Sun is practically parallel to the horizon. Appearing on the day of the vernal equinox, the Sun rises higher and higher for a quarter of a year, making circles above the horizon. On the day of the summer solstice, it reaches its maximum height (23.5˚). The next quarter of the year, before the autumnal equinox, the Sun descends. This is a polar day. Then the polar night sets in for six months. In mid-latitudes, throughout the year, the apparent diurnal path of the Sun either decreases or increases. It turns out to be smallest on the day of the winter solstice, and the largest on the day of the summer solstice. On the days of the equinox
The sun is at the celestial equator. At the same time, it rises at the point of the east and sets at the point of the west.
In the period from the vernal equinox to the summer solstice, the place of sunrise shifts slightly from the point of sunrise to the left, to the north. And the place of entry moves away from the west point to the right, although also to the north. On the day of the summer solstice, the Sun appears in the northeast, and at noon it culminates at its highest altitude in a year. The sun sets in the northwest.
Then the places of sunrise and sunset are shifted back to the south. On the winter solstice, the Sun rises in the southeast, crosses the celestial meridian at its minimum height, and sets in the southwest. It should be borne in mind that due to refraction (that is, the refraction of light rays in the earth's atmosphere), the apparent height of the star is always greater than the true one.
Therefore, the sun rises earlier and sets later than it would have been in the absence of the atmosphere.
So, the diurnal path of the Sun is a small circle of the celestial sphere, parallel to the celestial equator. At the same time, during the year, the Sun moves relative to the celestial equator to the north, then to the south. The day and night parts of his journey are not the same. They are equal only on the days of the equinox, when the Sun is at the celestial equator.
The expression "the path of the sun among the stars" will seem strange to some. After all, the stars are not visible during the day. Therefore, it is not easy to notice that the Sun is slowly, by about 1˚ per day, moving among the stars from right to left. But you can see how the appearance of the starry sky changes throughout the year. All this is a consequence of the revolution of the Earth around the Sun.
The path of the apparent annual movement of the Sun against the background of stars is called the ecliptic (from the Greek "eclipse" - "eclipse"), and the period of revolution along the ecliptic is called a sidereal year. It is equal to 265 days 6 hours 9 minutes 10 seconds, or 365, 2564 average solar days.
The ecliptic and the celestial equator intersect at an angle of 23˚26 "at the vernal and autumnal equinox points. At the first of these points, the Sun usually occurs on March 21, when it passes from the southern hemisphere of the sky to the northern one. In the second, on September 23, at the transition of their northern hemisphere At the ecliptic farthest to the north, the Sun occurs on June 22 (summer solstice), and to the south on December 22 (winter solstice) .In a leap year, these dates are shifted by one day.
Of the four points of the ecliptic, the main point is the vernal equinox. It is from it that one of the celestial coordinates is counted - right ascension. It also serves to count sidereal time and the tropical year - the time interval between two successive passages of the center of the Sun through the vernal equinox. The tropical year determines the seasons on our planet.
Since the vernal equinox moves slowly among the stars due to the precession of the earth's axis, the duration of the tropical year is less than the duration of the stellar year. It is 365.2422 solar average days. About 2 thousand years ago, when Hipparchus compiled his star catalog (the first one that has come down to us in full), the vernal equinox was in the constellation Aries. By our time, it has moved almost 30˚, to the constellation Pisces, and the point of the autumnal equinox has moved from the constellation Libra to the constellation Virgo. But according to tradition, the points of equinox are designated by the same signs of the previous "equinox" constellations - Aries and Libra. The same happened with the solstice points: summer in the constellation Taurus is marked with the sign of Cancer, and winter in the constellation Sagittarius is marked with the sign of Capricorn.
And finally, the last thing that is associated with the apparent annual motion of the Sun. Half of the ecliptic from the vernal equinox to the autumn (from March 21 to September 23), the Sun passes in 186 days. The second half, from the autumnal and spring equinox, - in 179 days (180 in a leap year). But the halves of the ecliptic are equal: each 180˚. Consequently, the Sun moves unevenly along the ecliptic. This unevenness is explained by a change in the speed of the Earth in an elliptical orbit around the Sun. The uneven movement of the Sun along the ecliptic leads to different lengths of the seasons. For residents of the northern hemisphere, for example, spring and summer are six days longer than autumn and winter. The Earth on June 2-4 is located from the Sun 5 million kilometers longer than January 2-3, and moves in its orbit more slowly in accordance with Kepler's second law. In the summer, the Earth receives from
The sun is less warm, but summer in the Northern Hemisphere is longer than winter. Therefore, the northern hemisphere of the Earth is warmer than the southern one.
Solar eclipses
At the time of the lunar new moon, a solar eclipse can occur - after all, it is on the new moon that the Moon passes between the Sun and the Earth. Astronomers know in advance when and where a solar eclipse will occur, and report this in astronomical calendars.
The Earth got a single satellite, but what a! The Moon is 400 times smaller than the Sun and just 400 times closer to the Earth, so in the sky the Sun and the Moon appear to be disks of the same size. So in a total solar eclipse, the Moon completely obscures the bright surface of the Sun, while leaving the entire solar atmosphere open.
Exactly at the appointed hour and minute, through the dark glass, one can see how something black creeps onto the bright disk of the Sun from the right edge, as a black hole appears on it. It gradually grows, until finally the solar circle takes the form of a narrow crescent. At the same time, daylight weakens quickly. Here the Sun completely hides behind a dark shutter, the last day's ray is extinguished, and the darkness, which seems deeper the more suddenly it is, spreads around, plunging man and all nature into silent surprise.
The English astronomer Francis Bailey tells about the eclipse of the Sun on July 8, 1842 in the city of Pavia (Italy): “When the total eclipse came and the sunlight instantly extinguished, around the dark body of the Moon suddenly appeared some kind of bright glow, like a crown or a halo around the head saint.
No reports of past eclipses have written anything like this, and I did not expect to see the splendor that was now before my eyes. The width of the corona, counting from the circumference of the Moon's disk, was approximately half the lunar diameter. She seemed to be composed of bright rays. Its light was denser near the very edge of the moon, and with the distance the rays of the corona became weaker and thinner. The attenuation of the light went perfectly smoothly along with the increase in distance. The crown was presented in the form of beams of direct weak rays; their outer ends fanned out; the rays were of unequal length. The crown was not reddish, not pearl, it was completely white. Its rays shimmered or flickered like a gas flame. No matter how brilliant this phenomenon was, no matter what delight it aroused in the audience, there was definitely something ominous in this strange, wondrous spectacle, and I fully understand how people could be shocked and frightened at a time when these phenomena happened completely unexpectedly.
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In 1802, the English physicist William Haid Wollaston (1766-1828), who discovered ultraviolet rays a year earlier, built a spectroscope in which a narrow slit was located in front of a glass prism parallel to its edge. Pointing the device at the Sun, he noticed that narrow dark lines intersect the solar spectrum.
Wollaston then did not understand the meaning of his discovery and did not attach special importance to it. 12 years later, in 1814. German physicist Joseph Fraunhofer (1787-1826) again discovered dark lines in the solar spectrum, but unlike Wollaston, he was able to correctly explain them by the absorption of rays by the gases of the Sun's atmosphere. Using the phenomenon of light diffraction, he measured the wavelengths of the observed lines, which have since been called Fraunhofer.
In 1833 g. Scottish physicist David Brewster (1781-1868), known for his studies of the polarization of light, drew attention to a group of bands in the solar spectrum, the intensity of which increased as the sun descended towards the horizon. Almost 30 years passed before, in 1862, the eminent French astrophysicist Pierre Jules César Jansen (1824-1907) gave them the correct explanation: these bands, called telluric bands (from Latin telluris - "earth"), are caused by the absorption of solar rays by gases earth's atmosphere.
By the middle of the XIX century. physicists have already studied the spectra of luminous gases quite well. So, it was found that the glow of sodium vapor generates a bright yellow line. However, at the same place in the spectrum of the Sun, a dark line was observed. What does that mean?
Solve this issue in 1859. undertook the outstanding German physicist Gustav Kirchhoff (1824-1887) and his colleague, the famous chemist Robert Boonsen (1811-1899). By comparing the wavelengths of Fraunhofer lines in the spectrum of the Sun and emission lines of vapors of various substances, Kirchhoff and Bunsen discovered sodium, iron, magnesium, calcium, chromium and other metals on the Sun. Each time, the glowing laboratory lines of terrestrial gases were matched by dark lines in the sun's spectrum. In 1862, the Swedish physicist and astronomer Andrei Jonas Angström (1814-1874), another of the founders of spectroscopy (by the way, the unit of length is named after him, angstroms: 1 A \u003d 10 ~ 10 m), discovered in the solar spectrum the lines of the most widespread in the nature of the element - hydrogen. In 1869 he, having measured the wavelengths of several thousand lines with great accuracy, compiled the first detailed atlas of the solar spectrum.
18 August 1868 French astrophysicist Pierre Jansen, observing a total solar eclipse, noticed a bright yellow line in the solar spectrum near the double sodium line. It was attributed to the chemical element helium, unknown on Earth (from the Greek "helios" - "sun"). Indeed, on Earth, helium was first found in gases released when the mineral cleveite was heated, only in 1895, so it fully justified its "extraterrestrial" name.
The advances in solar spectroscopy have stimulated scientists to apply spectral analysis to the study of the stars. An outstanding role in the development of stellar spectroscopy rightfully belongs to the Italian astrophysicist Angelo Sokchi (1818-1878). In 1863-1868. he studied the spectra of 4 thousand stars and built the first classification of stellar spectra, dividing them into four classes. Its classification was accepted by all astronomers and was applied until its introduction at the beginning of the 20th century. Harvard classification. Simultaneously with William Huggins, Secchi performed the first spectral observations of the planets, and he discovered in the red part of the spectrum of Jupiter a wide dark band, which, as it turned out later, belonged to methane.
A significant contribution to the development of astro-spectroscopy was made by compatriot Sekki Giovanni Donati(1826-1873), whose name is usually associated with the bright and very beautiful comet he discovered in 1858 and named after him. Donati was the first to obtain its spectrum and identify the bands and lines observed in it. He studied the spectra of the Sun, stars, solar chromosphere and corona, as well as auroras.
William Huggins (1824-1910) established the similarity of the spectra of many stars with the spectrum of the Sun. He showed that light is emitted by its incandescent surface, after which it is absorbed by the gases of the solar atmosphere. It became clear why the lines of elements in the spectrum of the Sun and stars are usually dark, not bright. Huggins was the first to obtain and study the spectra of gaseous nebulae, consisting of separate emission lines. This proved that they are gas.
Huggins first studied the spectrum of a new star, namely the new Northern Corona, which erupted in 1866, and discovered the existence of an expanding shell of gas around the star. He was one of the first to use the Doppler-Fizeau principle (often called the Doppler effect) to determine the velocities of stars along the line of sight.
Not long before that, in 1842, the Austrian physicist Christian Doppler (1803-1853) theoretically proved that the frequency of sound and light vibrations perceived by an observer depends on the speed of approach or removal of their source. The pitch of the horn of a locomotive, for example, changes sharply (downward) when an approaching train passes us and begins to recede.
The outstanding French physicist Armand Hippolyte Louis Fizeau (1819-1896) in 1848 tested this phenomenon for light rays in the laboratory. He also suggested using it to determine the velocities of stars along the line of sight, the so-called line-of-sight velocities, by the shift of spectral lines to the violet end of the spectrum (in the case of a source approaching) or to the red (in the case of its receding). In 1868 Huggins measured the radial velocity of Sirius in this way. It turned out that it is approaching the Earth at a speed of about 8 km / s.
The consistent application of the Doppler-Fizeau principle in astronomy has led to a number of remarkable discoveries. In 1889, the director of the Harvard Observatory (USA) Edward Charles Pickering (1846-1919) discovered the bifurcation of lines in the spectrum of Mi-tsar, a well-known 2nd magnitude star in the tail of the Big Dipper. Lines with a certain period moved, then moved apart. Pickering realized that this is most likely a close binary system: its stars are so close to each other that they cannot be distinguished in any telescope. However spectral analysis lets you do it. Since the velocities of both stars of the pair are directed in different directions, they can be determined using the Doppler-Fizeau principle (and, of course, the orbital period of the stars in the system).
In 1900 Pulkovo astronomer Aristarkh Apollonovich Belopolsky (1854-1934) used this principle to determine the speeds and periods of rotation of the planets. If we put the spectrograph slit along the planet's equator, the spectral lines will be tilted (one edge of the planet is approaching us, and the other is receding). Applying this method to the rings of Saturn, Belopolsky proved that parts of the ring revolve around the planet according to Kepler's laws, which means that they consist of many separate, unconnected small particles, as it was assumed, based on theoretical considerations, James Clerk Maxwell (1831- 1879) and Sofya Vasilievna Kovalevskaya (1850-1891).
Simultaneously with Belopolsky, the American astronomer James Edouard Kyler (1857-1900) and the French astronomer Henri Delandre (1853-1948) obtained the same result.
About a year before these studies, Belopolsky discovered a periodic change in the radial velocities in Cepheids. At the same time, the Moscow physicist Nikolai Alekseevich Umov (1846-1915) expressed the idea, ahead of his time, that in this case, scientists are not dealing with a binary system, as was then believed, but with the pulsation of a star.
Meanwhile, astrospectroscopy made more and more advances. In 1890, the Harvard Astronomical Observatory released a large catalog of stellar spectra containing 10,350 stars up to magnitude 8 and 25? south declination. It was dedicated to the memory of Henry Draper (1837-1882), an American amateur astronomer (specializing as a doctor), a pioneer in the widespread use of photography in astronomy. In 1872, he obtained the first photograph of the spectrum of a star (spectrogram), and later - the spectra of bright stars, the Moon, planets, comets and nebulae. After the release of the first volume of the catalog, additions to it were published more than once. The total number of studied spectra of stars reached 350 thousand.
Spectral analysis is the main method for determining the chemical composition of distant luminous objects, such as stars. The first elements discovered through this method were cesium and rubidium. And soon helium was also discovered, moreover, it was discovered on the Sun 27 years earlier than on Earth.
Everyone knows the seven primary colors recognized by our eye, but there are still shades in the transition from one color to another. Light is a mixture of electromagnetic waves, and each vibration has its own wavelength, and, accordingly, its own color. Passing light from an object through a prism, it is decomposed into spectra. From the resulting picture (spectrogram), conclusions are drawn about the characteristics of the object that emitted light. A real life example is a rainbow after a rain. Raindrops decompose the light flying from the sun into seven primary colors. Wavelength unit - Angstrem one hundred millionth part of a centimeter
All spectra that can be observed are divided into three classes:
- Linear emission spectrum. Emission lines are emitted by heated, rarefied gas.
- Continuous spectrum. These types of spectra are obtained from solids, liquids, and hot opaque gases.
- Linear absorption spectrum. The spectrum is formed if radiation from a hot body with a continuous spectrum passes through a rarefied cold medium.
Applications in astronomy
Spectral analysis is widely used in modern astronomy. This is a method capable of producing the most detailed and unique information about objects in space.
By analyzing the radiation of an object, you can very accurately establish its main characteristics.
The propagation of light is in the form of electromagnetic waves. Each color has a specific wavelength. The wavelength decreases in the spectrum from 7000 Angstroms to 4000 Angstroms, from red rays to violet. After the violet rays, the ultraviolet rays are located. They are not caught by the eye, but fixed by instruments. After ultraviolet, X-rays come - they have an even shorter wavelength.
The other side of the spectrum, red, continues with infrared rays, also invisible to the human eye, but captured by specially prepared photographic plates. Spectral observations are studies of rays in a range of colors from ultraviolet to infrared. The saturation of spectral lines determines the number of molecules and atoms that emit or absorb energy. The number of atoms is the greater, the brighter the line in the emitted spectrum and darker in the absorbed one. For all other stars, the presence is characteristic. Radiation passing through the atmosphere appears as dark absorption lines on the continuous spectrum of the visible surface. For such objects, these are absorption spectra. Spectral analysis, based on, allows you to determine the speed of movement of celestial bodies relative to our planet along the line of sight. At a source of light approaching the observer, the wavelengths are shortened, and if the source is moving away, then the wavelengths will increase. If a body moves on Earth, then its speed is negligible in the spectrum. And even the velocities of celestial bodies, having values \u200b\u200bof tens and hundreds of kilometers per second, are visible at such small displacements that their observation on spectrograms is really only possible with the help of a microscope. The resulting spectrogram of the luminary is compared with the standards, which are the spectrograms of terrestrial radiation sources, for example, a neon lamp. The shift of the spectral lines of the observed object is determined with respect to the stationary spectrum in the standards. This shift is very small, and its magnitude is calculated in tenths and hundredths of a millimeter.
Significance for cosmology
At present, all spectra of chemical elements have been determined and tabulated. Spectral the analysis revealed some unknown elements, such as rubidium and cesium. And these new elements sometimes received names corresponding to the colors of the dominant lines of the spectrum: rubidium gives dark red lines, and cesium (sky blue) - blue. Only spectral analysis helped to determine the chemical composition of our star and other stars.Using other methods to achieve this goal is not possible. As it turned out, the same chemical elements are present both on our planet and on distant stars. Astrophysics uses spectral analysis to recognize the characteristics of stars, gas clouds, and other objects. These are the chemical composition, temperature, speed of movement, magnetic induction, pressure. All these quantities are determined only by analyzing the spectral lines of space objects. By adopting the Doppler effect, it became possible to measure the radial velocities of thousands of stars, gas nebulae, and other extragalactic objects. The patterns of movement of individual stars and rotation of stellar systems were determined. The masses of galaxies and star clusters were established. Using the effect discovered by the Dutch physicist Zeeman, it is possible to determine the parameters of cosmic magnetic fields. Strong magnetic fields split the spectrum lines. Such an effect is also created by an electric field, which can arise in a star for a short time (Stark effect).
Spectral analysis - the most powerful tool for studying space objects.
A device for obtaining a spectrum - a spectroscope consists of a collimator, a prism and a telescope (Fig.). A narrow slit is installed in the front part of the collimator facing the light source. A diverging beam of rays goes from it into the collimator tube. The slit is located at the main focus of the collimator objective, so that a parallel beam of rays emerges from the collimator.
What happens if we direct this beam of rays into the objective of the third component of the spectroscope - the telescope?
Its lens will collect the rays at its main focus and an image of the slit is formed here; we can look at it through the eyepiece and see a clear image of the spectroscope entrance slit.
A triangular glass prism is placed between the objectives of the collimator and the telescope so that its refracting edge is parallel to the slit. The prism refracts the parallel beam of rays falling on it from the collimator lens, deflecting it towards its base. In this case, rays of different colors are deflected in different ways, depending on the wavelength, as follows from formula (3.2). Thus, the prism decomposes light into a set of monochromatic (monochromatic) beams of rays. Instead of one image of the slit in the focal plane of the telescope of the spectroscope, many multi-colored images of the slit are formed, adjacent to each other and distributed in accordance with the change in wavelengths, i.e., a rainbow band of the spectrum. The direction in which the spectrum is stretched is called the direction of dispersion. It is clear why the slit of the spectroscope should be narrow enough. If we widen the gap, then the adjacent monochromatic images will overlap and the spectrum will be "washed out".
During visual observations through a spectroscope, we see a rainbow strip of the spectrum. If, instead of an eyepiece, a cassette is placed in the focal plane of the telescope, then the telescope will turn into a photographic camera, and the spectroscope into a spectrograph - a device widely used by astrophysicists. True, with its help a black-and-white image of the spectrum is obtained, but this does not in the least interfere with obtaining the richest information about the heavenly bodies.
Figure: Spectroscope device
The spectrum of radiation emitted by an incandescent solid or a liquid heated to glow is continuous. If you look through the spectroscope at the filament of an electric light bulb, you can see a bright rainbow strip, which is called a continuous spectrum. There are methods that make it possible to measure the intensity of radiation at different wavelengths. Then, putting on the horizontal axis the wavelength I, and on the vertical axis - the radiation intensity (energy) E \\, we get a graph, which is called the energy distribution curve in the spectrum (Fig. 74). The appearance of this curve depends mainly on the temperature of the emitter. For rays with a short wavelength, the energy Eλ is small. As the wavelength increases, the energy increases and at a certain wavelength λ max reaches a maximum; with a further increase in the wavelength, the radiation energy decreases. It turns out that the temperature T and λ, max are related to each other by the formula
T x λ Max \u003d constant value.
This formula expresses Wien's law ( The formula includes the absolute temperature G, measured from the temperature t \u003d -273 ° Celsius.) It follows from it that little heated bodies emit long-wave (infrared) rays, while strongly heated bodies emit blue and even violet rays the most. By studying the distribution of energy in the spectrum, it is possible to determine the temperature of stars. This is one of the challenges posed by astrospectroscopy.
However, spectral studies make it possible to obtain much richer information about celestial bodies. The fact is that a rarefied gas heated to glowing does not emit a continuous spectrum, but a linear one, consisting of a certain set of narrow, almost monochromatic spectral lines. Bright lines are called emission lines. So, for example, if you introduce ordinary table salt into the flame of the burner, then it will turn into an intense yellow color. Through the spectroscope, we will see two bright yellow emission spectral lines, denoted D 1 and D 2, emitted by heated sodium vapor, which is part of table salt. The spectrum of iron converted at high temperatures into a gaseous state is especially rich in lines.
Detailed atlases and catalogs of spectral lines of chemical elements have been compiled, and this helps to perform spectral analysis of a substance, to find out what chemical elements are present in it.
It should be borne in mind that in addition to the emission lines, absorption, dark absorption lines are also observed, which occupy the same places in the spectrum. They are easy to observe in the laboratory if you do this experiment. Observing the continuous spectrum of an incandescent solid through the spectroscope, we place in the path of the rays, between this body and the slit of the spectroscope, the flame of a burner saturated with sodium vapor. In place of the two bright yellow sodium emission lines, we will see two dark lines D 1 and D 2 against the background of the continuous spectrum, since vapors and gases are able to absorb the same radiation that they themselves emit.
The form of the line spectrum, the wavelengths of the spectral lines depend on the properties of the given atom. As you know, the atom of any chemical element consists of a central, positively charged nucleus surrounded by electrons. The electron that is least strongly bound to the nucleus is more susceptible to external influences - it is called an optical electron. This electron is capable of absorbing radiation energy incident on the atom from the outside; "Stocking up" additional energy, it changes its movement, coming into an excited state. It can also come to an excited state as a result of collisions of an atom with another atom or electron, which are inevitable during thermal motion.
Atomic physics has established that each atom has its own certain discrete energy levels, and the electron, during its transitions, can "linger" only on them. Each of the levels can be assigned a certain number - the main quantum number; the higher this level is, the more its energy. Let us denote the energy corresponding to the quantum number k through E k, and the quantum number i - through E i and assume that E k is greater than E i. Let, further, an optical electron be excited to the state E k. According to the laws of atomic physics, an electron cannot remain in an excited state for a long time (with the exception of some energy levels) and after a millionth of a second, spontaneously, as they say, spontaneously, it must pass into another state with less energy.
Let us assume that it has also passed the state with energy E i. This transition is accompanied by the emission of a photon, the energy of which is equal to the difference ek - ei. The photon will have a frequency vfti, which is calculated by the formula
Hν ki \u003d E k - E i (3.5)
Where h is Planck's constant, equal to 6.6 X 10-27 erg "x sec. The photon has not only frequency, but also a wavelength λ \u003d c: ν, where c denotes the speed of light.
Thus, as a result of this transition, the optical electron will emit a discrete spectral line with a wavelength λ ki. Thus, a line emission spectrum is formed from various transitions of an optical electron.
In the normal, unexcited state, the electron has the energy of the deepest level, which we denote by E ±. Now let us assume that radiation of the most varied frequencies v falls on the atom from the outside. Can an optical electron absorb radiation of any frequency, i.e., any wavelength? Of course not, and here's why.
This atom has the following "allowed" energy levels, which we write out in ascending order:
Е 1, Е 2, Е 3, ... Е i, ..., Е k, ..., Е ∞
An electron can absorb radiation only of those frequencies that correspond to transitions
E 2 - E 1 \u003d hν 21, E 3 - E 1 \u003d hν 31, E 4 - E 1 \u003d hν 41, etc.
All these transitions correspond to discrete spectral lines with wavelengths
λ 21, λ 31, λ 41, etc.,
All of which together form a series of spectral lines corresponding to the absorption of radiation by an electron when it passes from the same energy level E 1.
If, before absorbing the radiation energy, the optical electron was already excited and was, for example, in a state with energy E 2, then it can absorb portions of energy
E 3 - E 2 \u003d hν 32, E 4 - E 2 \u003d hν 42, E 5 - E 2 \u003d hν 52
That is, again a set of discrete frequencies (hence, discrete wavelengths), but this time from a different series, which has a lower energy level E2.
Summarizing what has been said, we note that there are infinitely many series of spectral lines for a given atom, since they can start from any of the energy levels. In practice, one has to encounter only a small number of series, because as the quantum number corresponding to the lowest energy of the level defining a given series increases, the entire series shifts to the infrared part of the spectrum, the farther the "initial" quantum number of the given series is.
But one should not think that one atom as a result of a single act of energy absorption by one electron can absorb all the radiation of corresponding wavelengths available to it. As a result of a single absorption event, only one spectral line is formed. However, if there are many atoms and they are placed in a field of radiation with the most diverse frequencies, then all absorption lines will appear in the continuous spectrum of this radiation, united by the series described above. At the same time, radiation with intermediate wavelengths cannot be absorbed, and for it the "cloud" of atoms is transparent. In order to clearly understand the systematics of the spectral lines of a given chemical element, the permissible energy levels characteristic of it are arranged in the form of a diagram. Such a scheme for hydrogen atoms is shown in Fig. The greater the energy reserve of the optical electron, the higher the level is. Therefore, transitions from lower levels to upper levels correspond to absorption events, i.e., the formation of an absorption line (i.e., an absorption line). In the case of transitions from top to bottom, emission of an emission spectral line occurs.
To the left of each level are marked the main quantum numbers - the numbers of levels 1, 2, 3, 4, 5 and 6. The next, higher levels should be numbered 7, 8, 9, etc. ad infinitum. As the quantum numbers grow, the levels approach each other, and the energy level marked with oo corresponds to an infinitely large quantum number. If an electron in an unexcited state E 1 absorbs the energy corresponding to this level, then it loses its connection with the atom and leaves it into space, and the atom is ionized, acquiring an excess electric charge. Let us calculate this energy and the frequency of the absorbed radiation ν∞ 1. Then by formula (3.5) hν∞ 1 \u003d Е ∞ - Е 1. The frequency ν∞ 1 is called the frequency of the "head" of the series. It corresponds to the wavelength ν∞ 1.
Now let the optical electron have already been excited to the E 2 state. Then for the ionization of the atom it is necessary that the electron absorbs the energy
Е ∞ - Е 2 \u003d hν∞ 2,
Which corresponds to the frequency ν∞ 2 and the wavelength λ∞ 2. This is the wavelength of the second series head. Thus, each of the series has its own head.
But the electron can also absorb more energy, that is, even harder radiation with a shorter wavelength.
Then he leaves the atom with a residual energy 1/2 mυ 2, which can be calculated by the formula
1/2 mυ 2 \u003d hν - (Е ∞ - Е i), (3.6)
Where E i denotes the energy of the level at which the electron was at the moment of absorption of the photon.
Thus, in addition to the line spectrum, a continuous spectrum is formed.
Hydrogen is one of the most common chemical elements in the Universe, and we will have to meet with its properties more than once in the future. Therefore, we will consider it in more detail.
While in the normal Er state, an optical electron can absorb radiation having a wavelength of 1216 angstroms ( Angstrem is a unit of length equal to 10 -8 cm. It is designated by the letter A.). A Lyman series absorption line is formed, called the Lyman-alpha (L α) line. The electron goes into an excited state corresponding to the energy level E 2.
When energy is absorbed E 3 - E 1, the electron goes to the third level: a line is formed with a wavelength of 1026 A and is called the Lyman-beta line (Lβ). Its wavelength is shorter than that of L α. In the transition from the first level to the fourth, the spectral line L γ, with a wavelength of 973 A is absorbed. the area of \u200b\u200bcontinuous absorption comes into its own on the side of shorter waves. When harder radiation is absorbed, the hydrogen atom is ionized.
Under terrestrial conditions, the Lyman series in the spectra of celestial bodies cannot be observed, since the short-wavelength part of the spectrum with wavelengths less than 3200 A is completely absorbed by the terrestrial atmosphere. Thus, the Lyman series can be observed in laboratories or outside the earth's atmosphere from satellites and orbital observatories. This is one of the tasks of extra-atmospheric astronomy.
Transitions of an electron from the second (excited) level to higher levels give rise to the famous Balmer series, which is not absorbed by the atmosphere. It is clearly seen in the spectra of many stars.
When an optical electron passes from the second level to the third, an H α absorption line is formed, located in the red region of the spectrum. The Нр absorption line is formed during the transition of an electron from the second level to the fourth; it has a shorter wavelength than H α. This is followed by H γ, H δ, etc. The entire Balmer series converges to its head, which has a wavelength of 3646 A. At shorter wavelengths, we again encounter a region of continuous absorption, leading to ionization of the atom. This time the electron leaves the atom from the second level, from the excited state.
When an electron transitions from the third level to higher levels, a series of Paschen-Bak spectral lines is formed, located in the infrared region of the spectrum.
So far, we have dealt with the atomic absorption spectrum. All of the above can be applied to emission spectra of radiation. If an electron is in the upper excited state with energy E k, then it can, as we said, emit a photon of frequency ν ki, returning to a lower energy level E i. A bright line will appear in the spectrum - an emission line. In this case, there is often a "exchange" of one photon for several, with lower frequencies. Let's give a concrete example. Let us assume that as a result of absorption of radiation, the optical electron of a hydrogen atom has passed from the normal level E 1 to a level with energy E 4. This corresponds to the absorption of the spectral line L γ. Thereafter, the excited optical electron can have four possibilities of spontaneous transition to lower energy levels:
1) transition from the fourth level to the first, at which the same spectral line L γ is emitted;
2) transition from the fourth level to the second, and then from the second to the first; two spectral lines H β and L α are emitted;
3) transition from the fourth level to the third, and then from the third to the first; two spectral lines are emitted: Pashen - Baka α and H β;
4) transition from the fourth level to the third, then from the third to the second, and then from the second to the first; three spectral lines of Paschen are emitted - Baka α, H α and L α.
This phenomenon is observed in space rarefied gas. Note that when one photon is divided into several, each of the resulting spectral lines has a longer wavelength compared to the absorbed one.
A more detailed study of the atomic spectra and the structure of the electron shells of atoms led to the conclusion that each energy level E k corresponding to the principal quantum number k, consists of several sublevels. They are characterized, in addition to the main quantum number, also by secondary quantum numbers and differ somewhat from each other in the amount of energy; they now have different energies grouping around E k. According to the laws of atomic physics, not all transitions between sublevels can be carried out, or, as they say, allowed. There are cases when an excited optical electron, after emission of an allowed spectral line, enters a sublevel from which there is no allowed exit towards deeper energy levels, and it gets stuck in this state for a long time. Then they say that the electron got to a superstable, metastable level.
However, the laws of atomic physics do not know absolute prohibitions. If a transition from a metastable level by radiation is forbidden, this does not mean that it cannot be realized. The point is that the residence time of an electron at a metastable level is much longer than at a normal level. If during this time no external cause (for example, a collision with another atom or additional absorption of a photon) takes the electron out of the metastable level, then it will return to the normal state, emitting a "forbidden" spectral line.
For such a transition to occur, the gas must be very rarefied and the external radiation must be sufficiently weak. This is done, for example, in planetary nebulae and in the solar corona.
The hydrogen atom's deepest energy level E 1 consists of two sublevels, which differ in two possible different directions of rotation of the electron around the axis. Although these levels differ little in energy, one of them is slightly higher and is metastable. A small difference in energy values \u200b\u200bleads, in accordance with formula (3.5), to the fact that in the case of emission of a forbidden line, its frequency should be small, and, consequently, the wavelength should be large. Indeed, a hydrogen atom, being in outer space, emits a "radio line" with a wavelength of 21 cm.
We now turn to the description of molecular absorption spectra. They consist of more or less wide bands located in the spectrum regions characteristic of a given molecule. Each of the bands consists of a very large number of spectral lines, which are so closely spaced to each other that they can be separated only with spectral instruments with huge dispersion.
Molecular spectra are well studied in terrestrial laboratories, and this makes it possible to judge by the shape of the spectrum about the chemical composition of the light-absorbing medium through which the radiation passes. Molecules form and become stable at relatively low temperatures, for example, in the shells of cold (red) stars and in the atmospheres of planets.
Now it is necessary to mention one more phenomenon on which many important conclusions of astrophysics are based. We are talking about the Doppler principle, according to which when a light source moves along the line of sight, the wavelengths of spectral lines change in proportion to the speed. If the normal (laboratory) wavelength of any spectral line is equal to A, 0, and the observed wavelength is λ, then the following formula is valid
In it, c denotes the speed of light, and through υ r - the radial velocity, which is equal to the projection of the spatial velocity onto the line of sight. If the light source moves away, the wavelengths increase, and if it approaches, the wavelengths decrease. As it is easy to see from formula (3. 7), the radial velocities of the receding light source are positive, and the approaching ones are negative.
Until now, we have talked mainly about laboratory studies of spectra. When studying the spectra of celestial bodies, some special conditions must be taken into account.
Stars, including the Sun, are huge accumulations of gaseous matter heated to high temperatures. In their outer parts, the density and pressure of the gas are low, but they rapidly increase as they go deeper into the bowels. The temperature rises rapidly. Suffice it to say that if the temperature in the outer layers of the Sun is close to six thousand degrees, then near its center it reaches several million degrees. The rate of thermal motion of the gas is so high here that collisions of atoms lead to their complete ionization. Calculations show that such a substance is not very transparent to radiation. As we rise into the outer layers, the opacity decreases and, finally, we encounter such a layer from which the radiation we observe comes to us. This layer is called the photosphere.
The photosphere emits thermal radiation with a continuous spectrum; it arises due to the chaotic thermal motions of charged particles - electrons and ions.
More rarefied and colder layers are located above the photosphere, in which radiation coming from the photosphere is absorbed. The absorption spectrum described above is formed here. Thus, by studying the chemical composition of stars from their spectra, we are investigating the composition of stellar atmospheres, but not stellar interiors.
In the same way, by studying the additional spectral lines that appear in the spectrum of a planet in comparison with the spectrum of the Sun, we study the chemical composition of its atmosphere.
In addition, one should not forget that the Earth's atmosphere partially absorbs certain spectral lines and bands, which are called telluric. The absorption by oxygen molecules and water vapor is especially strong.
Is the study of the spectra of celestial bodies accessible to the amateur astronomy?
Of course, many problems, such as the determination of radial velocities, requiring the use of very complex, powerful and expensive equipment, cannot be solved by an amateur. At the same time, some spectral observations can be performed using very modest, sometimes homemade, instruments.
Spectral studies of planets are distinguished by a large depth of information and serve primarily for the qualitative and quantitative study of the chemical composition of atmospheres.
Passing through the planet's atmosphere, sunlight is scattered over the entire spectrum and absorbed in selected frequencies, after which absorption lines or bands appear in the planet's spectrum, completely analogous to telluric lines formed in the earth's atmosphere. If the planet's atmosphere contains the same gases as the Earth's atmosphere, then the corresponding lines (stripe) will simply merge with the telluric ones and strengthen them. But such an increase is difficult to notice when the planet's atmosphere is small or poor in the studied gas. In this case, the Doppler shift of the planetary lines relative to the telluric ones comes to the rescue, provided that the time is chosen for observing the planet when it moves most rapidly relative to the Earth (for elongations and quadratures). Of course, this method requires a high dispersion of the spectral apparatus, very dry weather when trying to detect water vapor, and in general - observations from high mountains in order to weaken the telluric lines. Better yet, make observations with telescopes lifted into the stratosphere or even beyond the Earth's atmosphere. After the successful flights of the AMS series "Venus", "Mars", "Mariner", "Viking", which analyzed the atmospheres of Venus and Mars from close distances or by direct sounding of the atmosphere, the described method lost its significance.
Another thing is the analysis of planetary atmospheres for gases that are absent or poorly represented in the earth's atmosphere. Then a simple comparison of the spectrum of the planet with the solar spectrum (it is more convenient to photograph the spectrum of the Moon) immediately makes it possible to say whether there is a given gas in the planet's atmosphere. Thus, carbon dioxide was discovered in the atmosphere of Venus (Fig. 195), and then the same discovery was made on the spectrum of Mars. One glance at the spectra of the outer planets is enough to see there powerful absorption bands, which, when compared with laboratory sources, turn out to be bands of ammonia and methane (Fig. 196).
The strongest absorption bands of water vapor, carbon dioxide, nitrogen oxide, and other gases of interest to astrophysicists are located in the infrared region of the spectrum. Unfortunately, the entire near infrared region from 1 to 100 μm contains powerful absorption bands of water vapor, so that the Earth's atmosphere is transparent to solar and planetary radiation only in the intervals between these bands, and two such gaps are in the vicinity of 4.2 μm and from 14 up to 16 microns - filled with very strong stripes.
(click to view the scan)
That is why searches for gases of planetary atmospheres, on the one hand, are profitable to produce in infrared rays, and on the other hand, this benefit is limited.
Ultraviolet radiation from the Sun, in turn, is very strongly absorbed in the atmospheres of planets, but this absorption is continuous, associated with the dissociation of the corresponding molecules. So, the dissociation of the ozone molecule makes the earth's atmosphere opaque in the area. At shorter wavelengths, the dissociation of oxygen and nitrogen is switched on, their ionization actively inhibits radiation with a wavelength of less than 1000 A. Of course, the study of planetary atmospheres based on these phenomena is possible only from vehicles flying above the earth's atmosphere. But in the atmospheres of planets the presence of gases with active continuous absorption in the regions of the spectrum closer to the visible is possible and this can serve as a means for analyzing the planetary atmosphere (see, for example, about ultraviolet absorption in the spectrum of Venus on p. 500). Many gases have absorption bands in the radio frequency range as well. The planet's own radio emission, passing through the atmosphere, is absorbed at certain frequencies and this can be detected during observations with a radio spectrograph by comparing the radiation intensity at the frequency of the band and at a nearby place in the spectrum.
The quantitative analysis of the chemical composition of planetary atmospheres is fraught with difficulties. As in the analysis of stellar atmospheres, the measure of radiation absorption is the equivalent width W of the line (CPA 420), which is part of a band or a solitary one, i.e., the lack of light in a line, expressed in units of radiation from an adjacent segment of the continuous spectrum. Of course, the equivalent width is primarily a function of the number of absorbing molecules on the path of a light beam from the Sun through the atmosphere to the planet's surface and back - through the atmospheres of the planet and the Earth - to the terrestrial observer. But, in addition to this dependence, the equivalent linewidth depends on the total density of the planet's atmosphere, that is, on the content of other gases in it, and on the atomic-molecular parameters that determine the given spectral transition.
If we know these latter, then from the observation of several bands, strong and weak, it is possible to determine both the partial pressure of a given gas and the total pressure of the atmosphere on the planet's surface, even if it remains unknown which gas prevails in the composition of the atmosphere. Those absorption bands, which consist of numerous strong lines, so that they merge with a relatively low dispersion, usually used in the infrared region, make it possible to find the product of the content in the atmosphere of a given gas (in atm cm) by the total atmospheric pressure, while the weak lines into the composition of a low-power strip, only the content of a given gas can be determined. It would seem that from here it is easy to find the total atmospheric pressure or, more precisely, the elasticity of gases at the base of the atmosphere, expressed in dyne / cm2 or mm Hg as indicated by an aneroid barometer (not mercury!).
Unfortunately, the final results do not deserve full confidence due to the uncertainty of the theory, and therefore the more correct way is to simulate the atmosphere by spectrographing sunlight that has passed many times inside a long tube filled with the gas under study at different pressures and different plausible impurities - nitrogen. oxygen, argon, etc., which could be found in the atmosphere of the inner planet (by analogy with the Earth), or hydrogen, helium in the case of outer planets. This method has only one weak point - the impossibility of reproducing in a narrow tube all the conditions of light scattering, which are realized in real planetary atmospheres.
An example of such a determination of the thickness of the atmosphere will be found further on p. 498, 513. Usually the power of the planet's atmosphere in relation to one or another gas is expressed in atmcm, that is, equal to the height of a column of gas at normal atmospheric pressure and temperature 0 ° C. This value is obviously directly proportional to the number of gas molecules in the atmosphere. For comparison, we present the content of various gases in the earth's atmosphere expressed in the same units:
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