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Unformatted text preview: Extra office hours: - Wed. March 2'“, 5:00 'to 6:30pm NASA, NOAO, ESAand The Hubble Heritage Team (STSCI/ The Sun (Kulner, Ch. 6; see also Bennett el al. Ch. 14, Shu, Ch. 5) AST 203 (Spring 2011) The Sun What do we see when we look at the Sun? And what is this telling us? We see only the outermost layers. Careful observations of the Sun allow us to build up a picture of the interior. 2003i'10i‘28 05:24 UT AST 203 (Spring 2011) The Sun The Sun acts as a blackbody T = 5800 K. Departures in the optical suggest that it is not a uniform T. Energy flows out of the Sun—core must be hotter than surface. AST 203 (Spring 2011) solar spectral energy distribution 10’ IO" 10“ 10“ II)" 10-11: urn 10.1. 10—16 10-” 10"" 10" 10'“ 10" 10" l 10‘ A (cm) Figure 5.1. The solar spectral energy distribution. (VSVNIOHOS) (nus) The Sun We know more about the Sun than any other star. We already know the mass of the Sun Radius can be found from its distance and its angular size L knowing the luminosity and blackbody temperature: R : 47TUT4 From these, we can compute its average mass density, MG 2 x 1033 g 3 p (4/3)ng (4/3)7r(7 x 1010 cm)3 g/Cm As we will see, the Sun is considerably more dense at its center and falls off as we approach the surface. AST 203 (Spring 2011) core: energy produced via thermonuclear reactions radiative and convective layers: transport energy outward photosphere: lower part of solar atmosphere—surface we see (T ~ 5800 K). chromosphere: ~ 104 km thick layer above photosphere; characteristic red color corona: faint, tenuous outermost layer of the solar atmosphere AST 203 (Spring 2011) We will do things in a slightly different order than your text. We will cover the appearance of the outer layers of the Sun first, and then discuss how radiation moves through the mater that makes up the Sun. AST 203 (Spring 2011) The Sun (VSVNIOHOS) 2003.1'101‘28 06:24 UT AST 203 (Spring 2011) (Fe XII at 195 angstroms imaged by the EIT instrument on SOHO) AST 203 (Spring 2011) Properties of the Sun - Mass = 2.0x1033 9 (333,000 Earth masses) - Diameter = 1.4x1011 cm (109 Earth Diameters) - Average Density = (MassNolume) = 1.4 g / cm3 - Luminosity (i.e., total power output) = 4x1033 ergs/s - Surface Temperature = 5800 K - Rotation Period (at equator) = 25 days - Distance from Earth = 1 AU = 1.5 x 1013 cm - Composition = 70% H, 28% He, 2% metals (by mass) The Sun is an average star in almost every way AST 203 (Spring 2011) Photosphere The light we see comes from the photosphere. Atoms can emit and absorb photons at specific wavelengths, but what about the continuum? Dominant source is H _ ions (hydrogen atom with extra, weakly bound electron). Absorption results from H’ + 7 a H + e’ bound-free process. Critical wavelength is 1640 nm lnverse process gives emission: H + 6* —> H’ + y In both cases, the photon can take on a range of energies,. AsT 203 (spring 2011) Photosphere The interaction of photons with H— ions make the photosphere opaque to visible light. Gas above this is transparent. The photosphere is the layer we see. AsT 203 (spring 2011) Granules Look at the Sun up close: granulation— light areas surrounded by darker areas. ~1000 km across Change on timescales of seconds. Dark areas are only a little bit cooler than the bright areas. Convection zonejust below I i j ' (S) ' photosphere. Hot material from deeper layers is rising (and appear bright) and — coolermaterialisfalling. O AST 203 (Spring 2011) (from Bennett et al.) Granules Emno sanchez-Andrade Nuflo. Klaus Gemam Puschmann. Franz Kneer. Julian Blanca Rodriguez Ii Nazarel Bella Gonzalez http://www.astro.physik.uni-goettingen.de/~bruno/APOD/apod.html http://antwrp.gsfc.nasa.gov/apod/ap090405.html AST 203 (Spring 2011) Photosphere Base of the photosphere = depth we see when we look at the center of the Sun. Effective temperature of almost 5800 K. Spectral lines tell us even more information. (NOAO) Strong lines from less abundant elements (Na and Ca)—balance of ionization and excitation. I , i. i i , V- I l I ll l l I I ll [ii ill 1 l r l 1 l l I lllll l 11"111 II1 I' I't ‘ II ‘ l il‘l‘l I ll I‘ll l1‘ll‘ I i ll | | l II d I'll:I “I‘IIIIP ‘llrl:Ill ‘ rm u ’1 filll‘ll‘ll 1 ‘ ll l‘ vl l 11" l 1 l i l i I l AST 203 (Spring 2011) Spectral Lines Absorption lines are not sharp— I.) . J phySIcal processes broaden them. 3 J g - J ., . i ’ .3 F J :4 J i One of the most important is Doppler ‘J J I M J broadening. Ar ‘ " J J J i . J =’ i J J . Atoms in the photosphere have I J J J random motions due to their J ‘ 1 I i _ thermal energy. ‘_ J J J “my .- ' “FWQM a WW" "weir (from Bennett et al.) The line-of—sight velocity results in a Doppler shift. For an ideal monatomic gas, the thermal energy is Eth : ng AST 203 (Spring 2011) Spectral Lines Average kinetic energy of the particles = thermal energy 1 2 _§ §m<v >—2k:T We call this the root mean square (rms) velocity: 3kT 1/2 vrms : < > m Doppler shift —> broadening, AA. In reality, it will be higher, since some atoms are moving faster than the average. Our estimate gives: 1/2 AA 2 E C m AST 203 (Spring 2011) Spectral Li ' Resulting line profile > as 0.6 Line is deepest at rest x wavelength—smaller Doppler 5 shifts more likely than larger. 0.4 0.2 0.0 )1, Consider Ha at 5500 K (rest wavelength of 656.28 nm). AA _ 1(311)”2 A0 0 m _ 3 X 1010 cm 8’1 1 3- 1.38 x 10—16 erg K—1 5800 K 1/2 _5 1.67 X 10*24 g I 4 X 10 AA = 0.03 nm AST 203 (Spring 2011) How Does Rotation Affect Spectral Lines? This side of the Sun is receding—redshift >< observer This side of the Sun is approaching—blueshift Rotating Sun (viewed from above) Line is broaded---line of sight to center is unshifted. Red and blue shifts at edges, Amos (Spring 2011) broaded the line about the rest wavelength The Chromosphere Most of the light we see comes from the photosphere, it accounts for only 0.1% of the radius of the Sun. Moving out from the photosphere: Chromosphere. Usually we do not see the Chromosphere, since it is optically thin Just before and after an eclipse, we can see it. If we take a spectra during the eclipse, we see very strong Ha emission (not absorption), indicating a temperature of 15000 K Interestingly, as we move outward from the photopshere, the temperature begins to increase! AST 203 rep. ing 2011) The Chromosphere Looking at the chromosphere in Ha, we see large scale features called supergranules. (B!P9d!>l!N\) (hemmequ JElOS 1299 Big) AST 20315pllng 2011i The Corona Above the chromosphere, we encounter the corona—even hotter, more tenuous. Gas is so rarefied—very small optical depth Only seen during eclipses. Corona is ~ 2 million K—emits X-rays. Theoretical challenge: how is corona heated to such great temperatures? AST 203 (Spring 2011) The Corona Photons from the photosphere pass right through the corona without doing any heating (since it is so optically thin). Two mechanisms that are often discussed as ways to heat the corona are wave heating and magnetic reconnection. These topics are beyond the scope of this class. Both involve the presence of an active magnetic field in the Sun. AST 203 (Spring 2011) Sunspots (VSVN) Sunspots: regions that are slightly cooler than their surroundings (about 3800 K)—appear dark. he . Sunspots have a structure consisting " i. of the inner umbra surrounded by the ' lighter penumbra. AST 203 (Spring 2011) Sunspots 400 Years of Sunspot Observations Modern Minmum 250 h E 3 E s 2 Maunder , Minimum 3 y ‘ ‘ : r 2.? ' . . - . , . 50 N O O m (vipedmwvuv BuiuueM quols) 1600 1650 1700 1750 1800 1850 1900 1950 2000 Number of sunspots on the Sun varies regularly on 11 year cycle. First noticed by E. Walter Maunder in 1904. Looking through historical records, very little activity observed from 1645 to 1715—Maunder minimum. AST 203 (Spring 2011) DAILY SUNSPOT AREA AVERAGED OVER INDIVIDUAL SOLAR ROTATIONS MN SUNSFOT AREA IN EQUAL AREA LATITUDE STRIPS (“/n OF STRIP AREA) l> 0.0% l> 0.1% l’l> 1.0% 30N 190 30S 1870 18110 12190 1900 1910 1920 1930 1940 1950 19(10 1970 1980 1990 2000 2010 DATE AVERAGE DAILY SUNSPOT AREA (“/n 0F VISIBLE HEAIISPHERE) 0.4 03 . ill l l‘ M l .llll l . ll M .‘11 l1 1 . “l ‘ l l 1 l. 0.0 1870 1880 1890 1900 1910 1920 1930 1940 1950 1960 1970 1980 1990 2000 2010 DATE hmm/xrlunrnmxh mm mmymanmmmmmnwr usmxsm/Hm-uwn my“. Lattiude of the Sunspot vs. time: butterfly diagram. At the start of the 11 year cycle, sunspots form preferentially at higher lattitudes. AST 203 (Spring 2011) Magnetic Fields Magnetic field lines represent the directions in which compass needles would point. tines closer together indicate a stronger magnetic fie/0’. Charged particles fol/ow spiraling paths along magnetic field lines. cowvlgme 2mm Pearson Educallon lnc uublighmg as Pealsnn Aumsmwasley (from Bennett et al.) AST 203 (Spring 2011) Magnetic Fields and Sunspots Sunspots associated with regions of enhanced magnetic fields. (Etpedmwvuietsi) Magnetic fields formed by moving charges (currents). Bar magnets: alignment of orbit and spin of the electrons Magnetic fields trap gas-nu»... T F: 5,803 K Magnetic field lines connect the north \ and south pole—form closed loops. \\ if \ Ole In the Sun moving charges act as a O dynamo to create the magnetic field. surrounding plasma from :lrd‘ T a 5,800 K 1 / (from Bennett et al.) AST 203 (Spring 2011) o 2006 Pearson bducmicn, 1m, publishing as Addison Westey Sunspots Zeeman effect: splitting of atomic energy levels when the atom is placed in a magnetic field. Allows us to measure the strength of the Sun's magnetic fields. The splitting of a spectral line is correlated with the position of the sunspot in the image at right. (ac)qu AST 203 (Spring 2011) Sunspots George Ellery Hale: sunspots associated with strong magnetic fields (1908). Occur in pairs—one spot acting is north pole, other is south pole. During 11 year cycle, the same pole will always lead the other as the Sun rotates. In the next 11 year cycle, the opposite polarity will lead—this means that the Sun's magnetic field reverses every 11 years. So the sunspot cycle is actually a 22 year cycle. AST 203 (Spring 2011) (stint/Bum 'a ')l) Magnetic Field Over A Cycle Fri:tl:[:.\ui;u11ii§ fill-Ell! Hiker-ruling- Magnetic field strength from 1991-2001. AST 203 (Spring 2011) Differential Rotation (:lSNNHnV/OSN) Sunspots change alignment as they move across the Sun—spots closer to the solar equator move faster than those near the pole. The Sun is rotating differentially! It takes 25 days for material at the equator to complete 1 rotation, but 28 days for material at 40° to complete 1 rotation. AST 203 (Spring 2011) Differential Rotation If the Sun's magnetic field was generated in the core, we would expect it to be stable. Earth's magnetic field reverses ~ every million years. The Sun's magnetic field is generated close to the surface. The magnetic field is dragged with the Sun's differential rotation, and it will wind up. n . . , r K ‘ (from Bennett et al.) AsT 203 (Sprim (D .006 Ramon Edutalmn. Inc" pubimhmg mAddimn WleLy Other Solar Activity Other activity is associate with magnetic fields. Solar flares are energetic ejections of particles lasting tens of minutes to hours. These flares can be seen in Ha emission, but radiate across the entire EM spectrum. Not completely understood— ejected particles appear to trace out magnetic field lines, usually associated with sunspots. (aovui/VSVN) AST 203 lSpI ing 2011 l The Solar Wind Extending far from the Sun is the solar wind—a stream of particles moving outward from the Sun. Observations tell us that the Sun is losing mass at the rate of 10'14 Melyr. AST 203 (Spring 2011) The Solar Wind The solar wind particles hit the earth and travel along the magnetic field lines—this gives rise to the aurora. E. 'U m E: n: v The solar wind is also what forces the tail of a comet to always point away from the Sun. (0Hos) AST 203 (Spring 2011) ...
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