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Interpretations Translations of the Book
Collier's Encyclopedia
THE SUN is: Interpretation
Translation
The SUN is a star around which the Earth and other planets of the Solar system revolve.
The sun plays an exceptional role for humanity as the primary source of most types of energy.
Life in the form we know would be impossible if the Sun shone a little brighter or a little weaker.
The sun is a typical small star, of which there are billions.
But because of its proximity to us, only it allows astronomers to study in detail the physical structure of the star and the processes on its surface, which is practically unattainable with respect to other stars, even with the help of the most powerful telescopes.
Like other stars, the Sun is a hot ball of gas, mainly consisting of hydrogen compressed by its own gravity.
The energy radiated by the Sun is born deep in its bowels during thermonuclear reactions that convert hydrogen into helium.
Seeping out, this energy is radiated into space from the photosphere - a thin layer of the solar surface.
Above the photosphere is the outer atmosphere of the Sun - the corona, which extends over many radii of the Sun and merges with the interplanetary medium.
Since the gas in the corona is very rarefied, its glow is extremely weak.
Usually invisible against the background of a bright daytime sky, the corona becomes visible only at the moments of total solar eclipses.
The gas density monotonically decreases from the center of the Sun to its periphery, and the temperature reaching 16 million degrees in the center.
K, decreases to 5800 K in the photosphere, but then increases again to 2 million k.
K in the crown.
The transition layer between the photosphere and the corona, observed as a bright red rim at the moments of total solar eclipses, is called the chromosphere.
The Sun has an 11 year cycle of activity.
During this period, the number of sunspots (dark areas in the photosphere), flares (unexpected flashes in the chromosphere) and prominences (dense cold clouds of hydrogen condensing in the corona) increases and decreases again.
In this article, we will talk about the above mentioned areas and phenomena on the Sun.
After a brief description of the Sun as a star, we will discuss its internal structure, then the photosphere, chromosphere, flares, prominences and corona.
The sun is like a star.
The sun is located in one of the spiral arms of the Galaxy at a distance of more than half of the galactic radius from its center.
Together with the neighboring stars, the Sun orbits the center of the Galaxy with a period of about 240 million years.
The sun is a yellow dwarf of spectral class G2 V, belonging to the main sequence in the Hertzsprung Ressel diagram.
The main characteristics of the Sun are given in Table 1.
Note that although the Sun is gaseous up to the very center, its average density (1.4 g/cm3) exceeds the density of water, and in the center of the Sun it is much higher than even gold or platinum, which have a density of approx. 20 g/cm3.
The surface of the Sun at a temperature of 5800 K emits 6.5 kW/cm2.
The sun rotates around the axis in the direction of the general rotation of the planets.
But since the Sun is not a solid body, different regions of its photosphere rotate at different speeds: the rotation period at the equator is 25 days, and at latitude 75° - 31 days.
Table 1.
CHARACTERISTICS OF THE SUN THE INTERNAL STRUCTURE OF THE SUN Since we cannot directly observe the interior of the Sun, our knowledge of its structure is based on theoretical calculations.
Knowing the mass, radius and luminosity of the Sun from observations, to calculate its structure, it is necessary to make assumptions about the processes of energy generation, the mechanisms of its transfer from the core to the surface and about the chemical composition of matter.
Geological data indicate that the luminosity of the Sun has not changed significantly over the past few billion years.
What kind of energy source can support it for so long?
Ordinary chemical gorenje processes are not suitable for this.
Even gravitational compression, according to the calculations of Kelvin and Helmholtz, could only maintain the glow of the Sun for about 100 million years.
Bethe solved this problem in 1939: the source of the Sun's energy is the thermonuclear conversion of hydrogen into helium.
Since the efficiency of the thermonuclear process is very high, and the Sun is almost entirely composed of hydrogen, this completely solved the problem.
Two nuclear processes provide the luminosity of the Sun: the proton proton reaction and the carbon nitrogen cycle (see also STARS).
Proton The proton reaction leads to the formation of a helium nucleus from four hydrogen nuclei (protons) with the release of 4.3 H10-5 erg of energy in the form of gamma rays,two positrons and two neutrinos for each helium nucleus.
This reaction provides 90% of the Sun's luminosity.
It takes 1010 years for all the hydrogen in the Sun's core to turn into helium.
In 1968, p.
Davis and his colleagues began to measure the flow of neutrinos generated during thermonuclear reactions in the core of the Sun.
This was the first experimental verification of the theory of a solar energy source.
Neutrino interacts very weakly with matter, so it freely leaves the bowels of the Sun and reaches the Earth.
But for the same reason, it is extremely difficult to register it with devices.
Despite the improvement of the equipment and the refinement of the model of the Sun, the observed neutrino flux still remains 3 times less than predicted.
There are several possible explanations: either the chemical composition of the Sun's core is not the same as that of its surface; or mathematical models of the processes occurring in the core are not quite accurate; or on the way from the Sun to the Earth, the neutrino changes its properties.
Further research in this area is needed.
See also NEUTRINO ASTRONOMY.
Radiation plays the main role in the transfer of energy from the solar interior to the surface, convection is of secondary importance, and thermal conductivity is not important at all.
At a high temperature of the solar interior, the radiation is mainly represented by X rays with a wavelength of 2-10 .
Convection plays a significant role in the central region of the core and in the outer layer lying directly under the photosphere.
In 1962, the American physicist R. Leighton discovered that sections of the solar surface oscillate vertically with a period of approx. 5 minutes.
Calculations by R. Ulrich and K. Wolf showed that sound waves excited by turbulent gas movements in the convective zone lying under the photosphere can manifest themselves in this way.
In it, as in an organ pipe, only those sounds are amplified, the wavelength of which exactly fits into the thickness of the zone.
In 1974, the German scientist F. Debner experimentally confirmed the calculations of Ulrich and Wolf.
Since then, the observation of 5 minute oscillations has become a powerful method for studying the internal structure of the Sun.
Analyzing them, it was possible to find out that: 1) the thickness of the convective zone is approx. 27% of the sun's radius; 2) the core of the Sun probably rotates faster than the surface;
3) the content of
the amount of helium inside the sun is approx. 40% by weight.
The observation of fluctuations with periods between 5 and 160 minutes was also reported.
These longer sound waves can penetrate deeper into the bowels of the Sun, which will help to understand the structure of the solar interior and, possibly, solve the problem of solar neutrino deficiency.
THE ATMOSPHERE OF THE SUN IS the photosphere.
This is a translucent layer several hundred kilometers thick, representing the" visible " surface of the Sun.
Since the atmosphere lying above is almost transparent, the radiation, having reached the photosphere from below, freely leaves it and goes into space.
Unable to absorb energy, the upper layers of the photosphere should be colder than the lower ones.
The proof of this can be seen in the photos of the Sun: in the center of the disk, where the thickness of the photosphere along the visual beam is minimal, it is brighter and bluer than at the edge (on the "limb") the disk.
In 1902, the calculations of A.Schuster, and later - E. Milne and A. Eddington confirmed that the temperature difference in the photosphere is just such as to ensure the transfer of radiation through a translucent gas from the lower layers to the upper ones.
The main substance that absorbs and re emits light in the photosphere is negative hydrogen ions (hydrogen atoms with an additional electron attached).
Fraunhofer spectrum.
Sunlight has a continuous spectrum with the absorption lines detected by Y. Fraunhofer in 1814; they show that in addition to hydrogen, many other chemical elements are present in the atmosphere of the Sun.
Absorption lines are formed in the spectrum because the atoms of the upper colder layers of the photosphere absorb light coming from below with certain wavelengths, and emit it not as intensively as the hot lower layers.
The distribution of brightness within the Fraunhofer line depends on the number and state of the atoms producing it, i.e. on the chemical composition, density and temperature of the gas.
Therefore, a detailed analysis of the Fraunhofer spectrum allows us to determine the conditions in the photosphere and its chemical composition (Table 2).
Table 2.
THE CHEMICAL COMPOSITION OF THE SUN'S PHOTOSPHERE Element is the Logarithm of the relative number of Hydrogen atoms_________12,00 Helium ___________11,20 Carbon __________8,56 Nitrogen _____________7,98 Oxygen _________9,00 Sodium ___________6,30 Magnesium ___________7,28 Aluminum _________6,21 Silicon __________7,60 Sulfur _____________7,17 Calcium __________6,38 Chrome _____________6,00 Iron ___________6,76 The most abundant element after hydrogen is helium, which gives only one line in the optical spectrum.
Therefore, the helium content in the photosphere is not measured very accurately, and it is judged by the spectra of the chromosphere.
No variations in the chemical composition in the Sun's atmosphere have been observed.
See also SPECTRUM.
Granulation.
Photos of the photosphere taken in white light under very good observation conditions show small bright dots - "granules" separated by dark gaps.
The diameters of the granules are about 1500 km.
They constantly arise and disappear, remaining 5-10 minutes.
Astronomers have long suspected that the granulation of the photosphere is associated with the convective movements of the gas heated from below.
Spectral measurements of J. Beckers proved that in the center of the pellet, hot gas really floats up at a speed of about 0.5 km / s; then it spreads out to the sides, cools down and slowly descends along the dark boundaries of the granules.
Supergranulation.
R. Leighton found that the photosphere is also divided into much larger cells with a diameter of approx. 30,000 km - "supergranules".
Supergranulation reflects the movements of matter in the convective zone under the photosphere.
In the center of the cell, the gas rises to the surface, spreads out to the sides at a speed of about 0.5 km / s and sinks down at its edges; each cell lives for about a day.
The movement of gas in supergranules constantly changes the structure of the magnetic field in the photosphere and chromosphere.
Photospheric gas is a good conductor of electricity (since some of its atoms are ionized), so the magnetic field lines appear to be frozen into it and are transferred by the movement of the gas to the boundaries of supergranules, where they concentrate and the field strength increases.
Sunspots.
In 1908, J. Hale discovered a strong magnetic field in the sunspots, coming out of the bowels to the surface.
Its magnetic induction is so great (up to several thousand gauss) that the ionized gas itself is forced to subordinate its movement to the configuration of the field; in the spots, the field slows down the convective mixing of the gas, which causes it to cool down.
Therefore, the gas in the spot is colder than the surrounding photospheric gas and looks darker.
Spots usually have a dark core - "shadow" - and a lighter "penumbra"surrounding it.
Usually, their temperature, respectively, is 1500 and 400 K lower than in the surrounding photosphere.
A SUNSPOT with a central shadow and a fibrous penumbra surrounding it.
Photospheric granules are visible around the spot.
The spot begins its growth from a small dark "pore" with a diameter of 1500 km.
Big most of the pores disappear after a day, but the spots that have grown out of them persist for weeks and reach a diameter of 30,000 km.
The details of the growth and decay of sunspots are not fully understood.
For example, it is not clear whether the magnetic tubes of the spot are compressed by the horizontal movement of gas or whether they are already ready to "emerge" from under the surface.
Harvey discovered in 1970 that the spots move towards the general rotation of the Sun faster than the surrounding photosphere (by about 140 m / s).
This indicates that the spots are associated with subphotospheric layers that rotate faster than the visible surface of the Sun.
Usually from 2 to 50 spots are combined into a group, often having a bipolar structure: at one end of the group there are spots of one magnetic polarity, and at the other - the opposite one.
But there are also multipolar groups.
The number of spots on the Sun's disk changes regularly with a period of approx. 11 years.
At the beginning of each cycle, new spots appear at high solar latitudes (± 50°).
As the cycle develops and the number of spots increases, they appear at lower and lower latitudes.
The end of the cycle is marked by the birth and decay of several spots near the equator (± 10°).
During the cycle, most of the "leading" (western) spots in the bipolar groups have the same magnetic polarity, and different in the northern and southern hemispheres of the Sun.
In the next cycle, the polarity of the leading spots changes to the opposite.
Therefore, they often talk about a full 22 year cycle of solar activity.
There is still a lot of mystery in the nature of this phenomenon.
Magnetic fields.
In the photosphere, a magnetic field with an induction of more than 50 Gs is observed only in spots, in the active regions surrounding the spots, as well as at the boundaries of supergranules.
But L. Stanflo and J. Harvey found indirect indications that the magnetic field of the photosphere is actually concentrated in thin tubes with a diameter of 100-200 km, where its induction is from 1000 to 2000 Gs.
Magnetically active regions differ from quiet regions only in the number of magnetic tubes per unit surface.
Probably, the solar magnetic field is generated in the depths of the convective zone, where the bubbling gas twists the weak initial field into powerful magnetic bundles.
The differential rotation of matter puts these bundles along parallels, and when the field in them becomes strong enough, they float up into the photosphere, breaking up in separate arches.
This is probably how spots are born, although there is still a lot of uncertainty about this.
The process of spot decay has been studied much more fully.
Supergranules that pop up at the edges of the active region capture the magnetic tubes and pull them apart.
Gradually, the general field weakens; an accidental connection of tubes of opposite polarity leads to their mutual destruction.
Chromosphere.
The chromosphere is located between the relatively cold, dense photosphere and the hot, sparse corona.
The weak light of the chromosphere is usually not visible against the background of a bright photosphere.
It can be seen as a narrow strip above the limb of the Sun, when the photosphere is closed naturally (at the time of a total solar eclipse) or artificially (in a special telescope coronograph).
The chromosphere can be studied throughout the entire disk of the Sun, if you observe in a narrow spectral range (about 0.5 ) near the center of a strong absorption line.
The method is based on the fact that the higher the absorption, the lower the depth to which our gaze penetrates into the Sun's atmosphere.
For such observations, a spectrograph of a special design is used - a spectrogeliograph.
Spectrogeliograms show that the chromosphere is heterogeneous: it is brighter above the sunspots and along the boundaries of supergranules.
Since it is in these areas that the magnetic field is strengthened, it is obvious that energy is transferred from the photosphere to the chromosphere with its help.
It is probably carried by sound waves excited by the turbulent movement of gas in the granules.
But in detail, the mechanisms of heating the chromosphere are not yet understood.
The chromosphere emits strongly in the hard ultraviolet range (500-2000), which is inaccessible to observation from the Earth's surface.
Since the early 1960s, many important measurements of the ultraviolet radiation of the upper atmosphere of the Sun have been made with the help of high altitude rockets and satellites.
More than 1000 radiation lines of various elements were found in its spectrum, including lines of repeatedly ionized carbon, nitrogen and oxygen, as well as the main series of hydrogen, helium and helium ion.
The study of these spectra showed that the transition from the chromosphere to the corona occurs at a distance of only 100 km, where the temperature increases from 50,000 to 2,000,000 K.
It turned out that the heating of the chromosphere largely occurs from the corona by thermal conductivity.
Near groups of sunspots in the chromosphere, bright and dark fibrous structures are observed, often elongated in the direction of the magnetic field.
Above 4000 km, uneven, jagged formations are visible, evolving quite quickly.
When observing the limb in the center of the first Balmer line of hydrogen (Ha), the chromosphere at these altitudes is filled with many spicules - thin and long clouds of hot gas.
Little is known about them.
The diameter of a single spicule is less than 1000 km; it lives for about 10 minutes.
With a speed of approx. 30 km/s spicules rise to a height of 10,000-15,000 km, after which they either dissolve or sink down.
Judging by the spectrum, the temperature of the spicules is 10,000-20,000 K, although the corona surrounding them at these altitudes is heated to at least 600,000 K.
It seems that spicules are areas of a relatively cold and dense chromosphere that temporarily rise into a hot, rarefied corona.
Counting within the boundaries of supergranules shows that the number of spicules at the level of the photosphere corresponds to the number of granules; there is probably a physical connection between them.
Flashes.
The chromosphere above a group of sunspots can suddenly become brighter and shoot a portion of gas.
This phenomenon, called "flash", is one of the most difficult to explain.
Flares emit powerfully in the entire range of electromagnetic waves from radio to X rays, and also often emit beams of electrons and protons at a relativistic speed (i.e. close to the speed of light).
They excite shock waves in the interplanetary medium that reach the Earth.
Flares are more likely to occur near groups of spots with a complex magnetic structure, especially when in
the group begins to rapidly grow a new spot; such groups produce several flashes per day.
Weak flashes occur more often than strong ones.
The most powerful flares occupy 0.1% of the solar disk and last for several hours.
The total energy of the flash is 1023-1025 J.
The X ray spectra of the flares obtained by the SMM (Solar Maximum Mission) satellite made it possible to significantly better understand the nature of the flares.
The beginning of the flash can be marked by an X ray burst with a photon wavelength of less than 0.05 , caused, as its spectrum shows, by a flow of relativistic electrons.
In a few seconds, these electrons heat up the surrounding gas to 20,000,000 K, and it becomes a source of X ray radiation in the range of 1-20 , hundreds of times higher than the flow in this range from a calm Sun.
At this temperature, iron atoms lose 24 of their 26 electrons.
Then the gas cools down, but still continues to emit X rays.
P. Wilde from Australia and A. Maxwell from the USA investigated the development of the flash using a radio analog spectrograph - a "dynamic spectrum analyzer" that registers changes in the power and frequency of radiation.
It turned out that the radiation frequency for the first few seconds of the flash drops from 600 to 100 MHz, indicating that a disturbance propagates through the corona at a speed of 1/3 of the speed of light.
In 1982, US radio astronomers using a VLA radio interferometer in pcs.
New Mexico and data from the SMM satellite, resolved fine details in the chromosphere and corona during the outbreak.
It is not surprising that these were loops, probably of a magnetic nature, in which energy is released that heats the gas during the flash.
At the final stage of the flare, relativistic electrons captured by the magnetic field continue to emit strongly polarized radio waves, moving in a spiral around the magnetic lines of force above the active region.
This radiation can last for several hours after the flash.
Although gas is always ejected from the flare area, its speed usually does not exceed the speed of departure from the surface of the Sun (616 km/s).
However, flares often emit streams of electrons and protons that reach the Earth in 1-3 days and cause auroras and magnetic field disturbances on it.
These particles with an energy reaching billions of electron volts are very dangerous for astronauts in orbit.
Therefore, astronomers try to predict solar flares by studying the configuration of the magnetic field in the chromosphere.
The complex structure of the field with twisted lines of force ready for reconnection indicates the possibility of a flash.
Prominences.
Solar prominences are relatively cold masses of gas that appear and disappear in a hot corona.
When observed with a coronagraph in the Ha line, they are visible on the limb of the Sun as bright clouds against a dark sky background.
But when observed with a spectrogeliograph or Lio interference filters, they look like dark fibers against the background of a bright chromosphere.
AN ERUPTIVE SOLAR PROMINENCE photographed during a total solar eclipse.
An eruptive (rising) prominence is formed from a dense cloud of gas ejected into space by the solar magnetic field.
The forms of prominences are extremely diverse, but there are several main types.
The prominences of sunspots are similar to curtains up to 100,000 km in length, 30,000 km in height and 5000 km thick.
Some prominences have a branched structure.
Rare and beautiful loop shaped prominences have a rounded shape with a diameter of approx. 50,000 km.
Almost all prominences have a thin structure of gas filaments, probably repeating the structure of the magnetic field; the true nature of this phenomenon is not clear.
The gas in the prominences usually moves in streams downwards at a speed of 1-20 km / s.
The exception is "sergi" - prominences that fly up from the surface at a speed of 100-200 km/s, and then fall back more slowly.
Prominences are born on the edges of groups of sunspots and can persist for several revolutions of the Sun (i.e., several Earth months).
The spectra of the prominences are similar to the spectra of the chromosphere: bright lines of hydrogen, helium and metals against the background of weak continuous radiation.
Usually, the radiation lines of quiet prominences are thinner than chromospheric lines; this is probably due to the smaller number of atoms on the visual beam in the prominence.
The analysis of the spectra indicates that the temperature of the quiet prominences is 10,000-20,000 K, and the density is about 1010 at.
/ cm3.
The active prominences show lines of ionized helium, which indicates a significantly higher temperature.
The temperature gradient in the prominences is very large, since they are surrounded by a corona with a temperature of 2,000,000 K.
The number of prominences and their distribution along the latitude during the 11 year cycle repeats the distribution of sunspots.
However, at high latitudes, there is a second belt of prominences, which shifts to the pole during the maximum of the cycle.
Why prominences are formed and what supports them in a sparse corona is not completely clear.
Crown.
The outer part of the Sun - the corona shines weakly and is visible to the naked eye only during total solar eclipses or with the help of a coronograph.
But it is much brighter in X rays and in the radio range.
See also EXTRA ATMOSPHERIC ASTRONOMY.
The corona shines brightly in the X ray range, because its temperature is from 1 to 5 million.
K, and at the moments of flashes reaches 10 million K.
X ray spectra of the corona have recently been obtained from satellites, and optical ones have been studied for many years during total eclipses.
These spectra contain lines of repeatedly ionized atoms of argon, calcium, iron, silicon and sulfur, which are formed only at temperatures above 1,000,000 K.
The SOLAR CORONA is the upper part of the Sun's atmosphere, which can be seen around the dark disk of the Moon at the time of a total solar eclipse.
The shape of the corona reflects the distribution of the magnetic field over the surface of the Sun.
The white light of the corona, which is visible up to a distance of 4 radii of the Sun during an eclipse, is formed as a result of scattering of photospheric radiation on free electrons of the corona.
Therefore, the change in the brightness of the corona with height indicates the distribution of electrons, and since the main element is completely ionized hydrogen, the distribution of the gas density also indicates.
Coronal structures are clearly divided into open (rays and polar brushes) and closed (loops and arches) the ionized gas exactly repeats the structure of the magnetic field in the corona, since it cannot move across the lines of force.
Since the field exits the photosphere and is associated with an 11 year sunspot cycle, the appearance of the corona changes during this cycle.
During the minimum period, the corona is dense and bright only in the equatorial belt, but as the cycle develops, coronal rays appear at higher latitudes, and at the maximum they can be seen at all latitudes.
From May 1973 to January 1974, the corona was continuously observed by 3 crews of astronauts from the Skylab orbital station.
