Chapter 9 Probing the Dynamic Sun

Solar Energy The Sun’s spectrum is close to that of an idealized blackbody with a temperature of 5800 K. Each square meter of the Sun’s surface emits a tremendous amount of radiation, principally at visible wavelengths. Because the Sun is so large, the total number of square meters of radiating surface is immense. The total amount of energy emitted by the Sun each second, called its luminosity, is about 3.9 x watts, or 3.9 x joules of energy emitted every second.

The Source of the Sun’s Energy The temperatures and pressures deep within the core of the Sun are so intense that hydrogen nuclei can combine to produce helium nuclei: thermonuclear fusion. In fusion, a small amount of matter is converted into MASSIVE amounts of energy.

Converting Hydrogen to Helium Each time this process takes place, a small fraction (0.7%) of the combined mass of the hydrogen nuclei does not show up in the mass of the helium nucleus. This “lost” mass is converted into energy. The Sun’s core contains enough hydrogen to continue doing so for more than 6 billion years into the future.

Hydrostatic and Thermal Equilibrium Equilibrium is maintained by a balance among three forces: 1. The downward pressure of the layers of solar material. 2. The upward pressure generated by hot gases. 3. The slab’s weight—that is, the downward gravitational pull it feels from the rest of the Sun.

Transporting Energy Outward from the Sun’s Core Thermonuclear reactions occur in the Sun’s core. Energy is transported outward, via radiative diffusion, to a distance of about 0.71 R. In the outer layers between 0.71 R and 1.00 R, energy flows outward by convection.

Modeling the Sun To construct a model of a star like the Sun, we express the ideas of hydrostatic equilibrium, thermal equilibrium, and energy transport as a set of equations. We then calculate conditions layer by layer in toward the star’s center. The result is a model of how temperature, pressure, and density increase with increasing depth below the star’s surface.

Probing the Sun’s Interior One powerful technique to infer what is going on beneath the Sun’s surface involves measuring vibrations of the Sun as a whole. This field of solar research is called helioseismology.

Observing the Photosphere The photosphere is the layer in the solar atmosphere from which the Sun’s visible light is emitted. The Sun appears darker around its limb, or edge. We see the upper photosphere, which is relatively cool and thus glows less brightly. The dark sunspots are also relatively cool regions.

Granulation High-resolution photographs of the Sun’s surface reveal a blotchy pattern called granulation. Granules are convection cells about 1000 km (600 mi) wide in the Sun’s photosphere. Rising hot gas produces bright granules. Cooler gas sinks downward along the boundaries between granules; this gas glows less brightly, giving the boundaries their dark appearance. This convective motion transports heat from the Sun’s interior outward to the solar atmosphere.

Supergranules and Large-Scale Convection Supergranules display little contrast between their center and edges, so they are hard to observe in ordinary images. In a false-color Doppler image like this one, light from gas that is approaching us (i.e., rising) is shifted toward shorter wavelengths, whereas light from receding gas (i.e., descending) is shifted toward longer wavelengths.

The Sun’s Chromosphere During a total solar eclipse, the Sun’s glowing chromosphere can be seen around the edge of the Moon. It appears pinkish because its hot gases emit light at certain wavelengths, principally the H a emission of hydrogen at a red wavelength of nm. Spicules are jets of chromospheric gas that surge upward into the Sun’s outer atmosphere.

The Corona The corona – is the outer layer of the Sun’s atmosphere, extending out to a distance of several million kilometers – is only about one-millionth as bright as the photosphere and can be viewed only when the light from the photosphere is blocked out – looks like numerous streamers extending in different directions far above the solar surface, changing over days and weeks – has temperatures far greater than the temperatures in the chromosphere

Temperatures in the Sun’s Upper Atmosphere The corona is actually not very “hot,” containing very little thermal energy. The corona is nearly a vacuum, but the atoms there are moving at very high speeds. But because there are so few atoms in the corona, the total amount of energy in these moving atoms (a measure of how “hot” the gas is) is rather low.

The Solar Wind Solar Wind is the outflow of coronal gases, traveling at a million kilometers per hour. Each second the Sun ejects about a million tons of material into the solar wind, composed almost entirely of electrons and nuclei of hydrogen and helium. The aurorae seen at far northern or southern latitudes on Earth are produced when electrons and ions from the solar wind enter our upper atmosphere.

Observing Sunspots A sunspot is a region in the photosphere where the temperature is relatively low, which makes it appear darker than its surroundings. Although sunspots vary greatly in size, typical ones measure a few tens of thousands of kilometers across—comparable to the diameter of the Earth. Sunspots are not permanent features of the photosphere but last between a few hours and a few months.

The Sunspot Cycle The average number of sunspots on the Sun is not constant, but varies in a predictable sunspot cycle. A period of exceptionally many sunspots is a sunspot maximum, as occurred in 1979, 1989, and 2000 and projected to occur in Conversely, the Sun is almost devoid of sunspots at a sunspot minimum, as occurred in 1976, 1986, 1996, and 2008.

Sunspots Are Associated with Magnetic Fields

The Magnetic-Dynamo Model Magnetic field lines tend to move along with the plasma in the Sun’s outer layers. Because the Sun rotates faster at the equator than near the poles, a field line that starts off running from the Sun’s north magnetic pole (N) to its south magnetic pole (S) ends up wrapped around the Sun. Sunspot groups appear where the concentrated magnetic field rises through the photosphere.

Magnetic Arches Plasma tends to follow the Sun’s magnetic field, with streamers of electrically charged particles moving along each field line. When the magnetic fields of two arches come into proximity, they can rearrange and combine. The tremendous amount of energy stored in the magnetic field is then released into the solar atmosphere.

Prominences Magnetic fields can also push upward from the Sun’s interior, compressing and heating a portion of the chromosphere that appears as bright, arching columns of gas called prominences. These can extend for tens of thousands of kilometers above the photosphere. Some prominences last for only a few hours, while others persist for many months. The most energetic prominences break free of the magnetic fields that confined them and burst into space.

Coronal Mass Ejections In a coronal mass ejection, more than a billion tons of high-temperature coronal gas is blasted into space at speeds of hundreds of kilometers per second. A typical coronal mass ejection lasts a few hours. These explosive events seem to be related to large-scale alterations in the Sun’s magnetic field.

The Aurorae Some solar wind particles are able to leak through Earth’s magnetic field at its weaker points and cascade down into the Earth’s upper atmosphere to where Earth’s magnetic field connects with the Earth near the planet’s north and south poles. As these high-speed charged particles collide with atoms in the upper atmosphere, they excite the atoms to high energy levels. The atoms then emit visible light as they drop down to their ground states, like the excited gas atoms in a neon light.