The Sun Our Star

1. The SunOur Star © Sierra College Astronomy Department

2. The Sun - Our StarSolar Data • As viewed from the Earth, the Sun has an average angular diameter of 31’ 59” and is at an average distance from the Earth of 1.50 X 108 km. • The diameter of the Sun can therefore be calculated as 1.39 X 106 km, or about 110 times Earth’s diameter and about 10 times Jupiter’s. © Sierra College Astronomy Department

3. The Sun - Our StarSolar Data • The mass of the Sun is 1.99 X 1030 kg, more than 300,000 times Earth’s mass. • The Sun’s average density is 1.41 g/cm3, about the same as the density of Jupiter. • The escape velocity of the Sun is 620 km/s. • The Sun rotates in 25.4 days at its equator and 36 days near its poles. © Sierra College Astronomy Department

4. The Sun - Our StarSolar Data • The Sun emits energy is all portions of the electromagnetic spectrum with its peak in the visible portion. • Luminosity is the total power output of the Sun, which is equal to 3.8 X 1026 watts. • Brightness: Solar energy strikes the Earth at the rate of 1,380 watts/m2 (Solar Constant). • We can examine the surface of the sun by looking at its spectra (e.g., the Fraunhofer lines). • The Sun is a G2V star with an absolute magnitude of 4.8. © Sierra College Astronomy Department

5. The Sun - Our StarBasic Solar Physics • To estimate the internal parameters of the Sun, astronomers must solve a set of equations that describe how pressure, temperature, mass, and luminosity change. • We need to know the chemical composition, rate of energy production and surface conditions to solve these equations. © Sierra College Astronomy Department

6. The Sun - Our StarSolar Energy Source Must produce 3.9 x 1026 Watts • Combustion • Hydrogen + Oxygen produces energy but would last only 1000 years • Meteorite Impacts • The infall of material can be converted into heat and light • About 1/10 Earth masses of debris would have to fall into the Sun each year • Gravitational contraction • The infall of material produces energy • Could sustain sun for about 500 million years • Still short of geologic age of Earth (~ 4.5 billion yr) © Sierra College Astronomy Department

7. The Sun - Our StarSolar Energy Source • Mass/Energy Conversion • Einstein proposed in 1905 that mass and energy are interconvertible - derives famous E = mc2 equation. • Nuclear Fission • Large atoms break into two and release energy • Problem: Sun does not contain enough of the big atoms • Nuclear Fusion • Light atoms combine to produce larger atom and release energy • Sun full of small atoms • Turns out to be the right answer © Sierra College Astronomy Department

8. The Sun - Our StarSolar Energy Source • In nuclear fusion, two nuclei (consisting of protons and neutrons) combine to form a larger nucleus, releasing energy in the process. • Large temperatures (15.6 million K) and densities (150,000 kg/m3) are needed to overcome electromagnetic repulsion between the protons. • This only occurs in the Sun’s core (inner 10% of its radius). • In the Sun,4hydrogen nuclei are fused to form1helium nucleus and energy (proton-proton chain). © Sierra College Astronomy Department

9. The Sun - Our StarSolar Energy Source • To produce the Sun’s energy output nearly 5 million tons of matter must be converted into energy each second. • This in turn requires 610 billion kg of hydrogen be transformed into 606 billion kg of helium. • Solar lifetime: ~ 10 billion years • But how do we confirm this from observations? © Sierra College Astronomy Department

10. The Sun - Our StarSolar Neutrinos • Neutrinos are highly non-interacting particles which can be used as a direct probe into the Sun (if our nuclear fusion idea is right) • However, they are incredibly hard to detect • Huge underground detectors are used to look for neutrinos from the Sun • Until recently there seemed to be too few coming from the Sun • Solution: neutrinos have a little mass and come in different “flavors” and change identity on the way to Earth! © Sierra College Astronomy Department

11. The Sun - Our StarBasic Interior Structure Hydrostatic Equilibrium • In a star or a planet, hydrostatic equilibrium is the balance between pressure caused by the weight of material above and the upward pressure exerted by material below. • Pressure at the Sun’s center is calculated to be 1.3 X 109 times that on the surface of the Earth. © Sierra College Astronomy Department

12. The Sun - Our StarBasic Interior Structure • Helioseismology studies the Sun’s interior by observing pulsations in the photosphere. • Many patterns or modes of oscillationcan be seen and resemble those from a string on a violin or seismic waves from earthquakes. • Typical speeds of up/down motions are 100 to 300 m/s with periods of 5 minutes. • Short wavelength modes of a few 1000 km resonate only in outermost parts of Sun. Longer wavelengths penetrate deeper. © Sierra College Astronomy Department

13. The Sun - Our StarBasic Interior Structure • Helioseismology has been used to detect: • The bottom of the convection zone • The Sun has more helium than previously thought • The interior has surprisingly about the same rotation period as the surface • An extensive gas stream system beneath the surface and as much as 30% of the way to the Sun’s center • Sunspots on the far side of the Sun © Sierra College Astronomy Department

14. The Sun - Our StarEnergy Transport • There are three possible mechanisms for energy transport from one location to another: conduction, radiation, and convection. • Conduction is the transfer of energy within a substance by collisions between atoms and/or molecules. This is not a significant factor in transporting energy within the Sun. © Sierra College Astronomy Department

15. The Sun - Our StarEnergy Transport Radiation • The energy produced by the Sun is carried out by photons emitted at one spot and absorbed at another (radiative diffusion). • The speed at which photons get out from the Sun depends heavily on its opacity– the ability of a substance to stop photons (the opposite of transparency). • In the Sun’s core, the photons typically travel 10-6 m before being reabsorbed. • About 1025 absorptions and reemissions are needed before the energy reaches the Sun’s surface. © Sierra College Astronomy Department

16. The Sun - Our StarEnergy Transport Convection • Radiative diffusion carries solar energy out to 70% of the Sun’s radius. • At this distance, the temperature has dropped to 1.5 million K and hydrogen atoms start to from, increasing the opacity. • As a result, the rate of temperature decline becomes steep and convection takes over the energy transport in the outer 30% of the Sun’s radius • Typically, the diffusion of energy from core to surface takes on average 170,000 years. © Sierra College Astronomy Department

17. The Sun - Our StarThe Photosphere • Photosphereis the visible “surface” of the Sun. • It is the part of the solar atmosphere from which light (most energy is in optical) is emitted into space. • The photosphere is a very thin layer - 200 km thick. • The photosphere varies in temperature from about 8,000 K at its deepest to 4,000 K near its outer edge. • Overall, the light received from the photosphere is representative of an object about 5,800 Kelvin. • Pressure of the outer photosphere is 0.01 the pressure of the Earth’s surface. The calculated density of particles is 0.001 of the density of air at sea level. © Sierra College Astronomy Department

18. The Sun - Our StarThe Photosphere • Limb is the apparent edge of the Sun (or any object) seen in the sky. • The edge of the Sun appears dimmer than the center. This is known aslimb darkening. • Observing the limb of the Sun one sees to a lesser depth because the line of sight is at a grazing angle. © Sierra College Astronomy Department

19. The Sun - Our StarThe Photosphere • Granulationis the division of the Sun’s surface into small convection cells. • Granules are areas where hot material (light areas) is rising from below and then descending (dark surroundings). • Each granule is typically 1000 km across, separated by regions 100 K cooler, and last about 15 minutes. • Supergranulation is convective pattern on a much larger scale, typically 30,000 km across, lasting about 1 day, and extending to the bottom of the convection zone. © Sierra College Astronomy Department

20. The Sun - Our StarThe Photosphere • Mass of the photosphere is 71% hydrogen, with helium comprising most of the remainder and a few percent consisting of several elements found on Earth. • From our knowledge of nuclear fusion, we know the Sun’s core must hold more helium. Calculations show that the hydrogen makes up only 34% of the center. © Sierra College Astronomy Department

21. The Sun - Our StarThe Lower Atmosphere • The chromosphere is the region of the solar atmosphere some 2,000 to 3,000 km thick that lies between the photosphere and the corona. • It is not usually observable from Earth except during a total solar eclipse. • When seen, it has a reddish color, which is caused by the hydrogen Balmer a emission line. © Sierra College Astronomy Department

22. The Sun - Our StarThe Lower Atmosphere • The chromosphere is mostly empty with the exception of spicules, which: • Are narrow jets of gas originating at the base of the chromosphere, shooting upward with speeds of about 25 km/sec, and lasting about 5 minutes • form a network above the edges of the photospheric supergranules © Sierra College Astronomy Department

23. The Sun - Our StarThe Lower Atmosphere • The corona is the outermost portion of the Sun’s atmosphere that can only be seen during a total solar eclipse. • Much of the light of the corona originates in the photosphere – light which is then scattered by electrons in the corona. • Some light is also the result of emission from ionized elements. • The temperature change from the chromosphere to the corona is quite dramatic rising from about 4000 K to 1 million K in a transition zone only a few thousand km thick. • Heating of the chromosphere and especially the corona is not understood and remains one of the unanswered questions about the Sun (although the heating process most likely involves energy transport from the convection zone via the Sun’s time-varying magnetic fields). © Sierra College Astronomy Department

24. The Sun - Our StarThe Lower Atmosphere Coronal Structure • Best seen in X-ray and ultraviolet images • Bright areas, also referred to as active regions, are closed magnetic field regions with hot gas often seen along coronal loops • Darker and cooler areas are coronal holes, which are open magnetic field regions emitting high-speed gas into interplanetary space © Sierra College Astronomy Department

25. The Sun - Our StarThe Lower Atmosphere Prominences • Dense, relatively cool regions reaching upwards of 50,000 km or more into corona and lasting 2-3 months • Appear as dark filaments against brighter disk of Sun. • Quiescent prominences located away active regions live longer than active prominences above active regions. • When a prominence reaches about 50,000 km it erupts sending gas into interplanetary space with speeds in excess of 1000 km/s. These eruptions are called coronal mass ejections. © Sierra College Astronomy Department

26. The Sun - Our StarThe Lower Atmosphere Flares • Solar flares, much more explosive and energetic than prominences, are erupt releases of magnetic energy. • Flares take place in active regions where a prominence is supported against gravity by magnetic field lines and then the magnetic field structure changes abruptly. • Coronal gas may heat to 40 million K and X-rays and ultraviolet light are emitted. The Sun’s brightness may increase by 1% during an unusually bright flare. • Flares blast out large numbers of very energetic charged particles that will take about 3 days to reach Earth. • Large solar flares cause spectacular auroras and can affect earthly radio transmissions if the ionosphere is disrupted by high-energy particles. © Sierra College Astronomy Department

27. The Sun - Our StarThe Solar Wind and Heliosphere • The Solar Wind • The flow of coronal gas into interplanetary space is called the solar wind. • Coronal gas takes about 4 days to reach the Earth. • Near the Earth the solar wind travels at 450 km/s; density is from 2–10 particles/cm3. • Speeds in excess of 1000 km/sec are possible during coronal mass ejections. • The Heliosphere • The boundary between the region dominated by the solar magnetic field and interstellar space is called the heliopause. • The region interior to the heliopause is the heliosphere. © Sierra College Astronomy Department

28. The Sun - Our StarSunspots and the Solar Activity Cycle • Dark spots on the Sun were first reported by the Chinese in the 5th century B.C. • Galileo and Thomas Harriott were the first Europeans to report these sunspots in the early 17th century. • In 1851, Schwabe discovered the sunspot cycle, which lasts about 11 years. • This periodic cycle, however, is not always present as the Maunder Minimum appears to show. • Additionally, individual sunspots are temporary phenomena lasting from a few hours to a few months. © Sierra College Astronomy Department

29. The Sun - Our StarSunspots and the Solar Activity Cycle • Magnetic fields can be measured using the Zeeman effect, which is the splitting of spectral lines by a strong magnetic field. Spectra No magnetic Field Magnetic Field • Sunspots involve the Sun’s magnetic field -The magnetic field in a sunspot is about 1,000 times that of the surrounding photosphere. © Sierra College Astronomy Department

30. The Sun - Our StarWhy are sunspots dark? • Sunspots are dark because they are cooler than there surrounding photosphere. • They are cooler because the magnetic field associated with the sunspot disrupts the convection in the granule • . • Heat does not flow as efficiently to the photosphere • Spot becomes cooler than normal (it is seen as dark, though it would be red if seen in isolation) • Sunspots, which often appear in pairs aligned in an east-west direction, have opposite magnetic polarities - one being north and the other south. © Sierra College Astronomy Department

31. The Sun - Our StarSunspots and the Solar Activity Cycle • Details of the Solar Sunspot Cycle • At a sunspot maximum, most spots occur about 30° north or south of the equator. • As the sunspot cycle progresses, the spots are seen closer and closer to the Sun’s equator. • When the spots reach the equator, the cycleis at asunspot minimumand begins again. • Location and relative number of sunspots can be plotted on a butterfly diagram. © Sierra College Astronomy Department

32. The Sun - Our StarSunspots and the Solar Activity Cycle A Model for the Sunspot Cycle • The current model (The Babcock model) has patterns of magnetic field lines within the Sun’s interior forming tubes that gradually become twisted from the Sun’sdifferential rotation. • Forced to the surface, these magnetic tubes become visible as sunspots. • Breaking the surface weakens the field lines and the sunspots die out. • As new lines form deep within the Sun, the magnetic field direction in the emerging tubes will reverse over time, causing the magnetic field of the Sun to reverse in direction in a 22-year period. This is called the magnetic cycle of the sun. © Sierra College Astronomy Department