Spectral Line Forma0on •  Astronomers can glean an enormous amount of informa1on about the Sun from an analysis of the absorp1on lines that arise in the photosphere and lower atmosphere. •  A detailed spectrum of the Sun spanning a range of wavelengths from 3600 to 6900 Å. No1ce the intricate dark Fraunhofer absorp1on lines superposed on the background con1nuous spectrum, which indicate the presence of some 67 different elements in various stages of excita1on and ioniza1on in the lower solar atmosphere. •  Spectral lines arise when electrons in atoms or ions make transi1ons between states of well-­‐defined energies, emiJng or absorbing photons of specific energies (that is, wavelengths or colors) in the process. •  Photons with energies well away from any atomic transi1on can escape from rela1vely deep in the photosphere, but those with energies close to a transi1on are more likely to be reabsorbed before escaping, so the ones we see on Earth tend to come from higher, cooler levels in the solar atmosphere. The inset shows a close-­‐up tracing of two of the thousands of solar absorp1on lines, the “H” and “K” lines of calcium at about 395 nm (3950 Å). •  As illustrated in the figure above, when we look at the Sun, we are actually peering down into the solar atmosphere to a depth that depends on the wavelength of the light under considera1on. •  Photons with wavelengths well away from any absorp1on feature tend to come from deep in the photosphere. However those at the centers of absorp1on lines, being much more likely to interact with maUer as they travel through the solar gas, mainly escape from higher, cooler levels. •  The lines are darker than their surroundings because the temperature at the level of the atmosphere where they form is lower than the 5800 K temperature at the base of the photosphere, where most of the con1nuous emission originates. (Recall that by Stefan’s law, the brightness of a radia1ng object depends on its temperature—the cooler the gas, the less energy it radiates.) •  The existence of the Fraunhofer lines is direct evidence that the temperature in the Sun’s atmosphere decreases with height above the photosphere. The Chromosphere and the Transi0on Zone •  The temperature of the solar atmosphere decreases with increasing radius through the photosphere and reaches a minimum (∼4200 K) in the lower chromosphere. •  It then rises to an intermediate plateau before increasing rapidly in a transi1on zone, eventually reaching millions of degrees in the solar corona. •  Different regions of the solar atmosphere can be studied by observa1ons at different wavelengths within spectral line profiles. Strong lines saturate quickly in the line center and thus represent the highest level of the solar atmosphere. Fig. IV.7 taken from “Physics and Chemistry of the Solar System”, John S. Lewis, p87 •  The change of gas temperature in the lower solar atmosphere is drama1c. •  The minimum temperature occurs in the chromosphere. Beyond that, the temperature rises sharply in the transi1on zone, finally leveling off at around 3 million K in the corona. •  The temperature of the Sun’s atmosphere varies with al1tude. The temperature decreases to a minimum of about 4500 K some 500 km above the photosphere, acer which it rises steadily. •  About 1500 km above the photosphere, the gas temperature begins to rise rapidly, reaching more than 1 million K at an al1tude of 10,000 km. Thereacer, in the corona, the temperature remains roughly constant at around 3 million K. •  Ultraviolet, x-­‐ray, and radio regions of the spectrum are used to explore the outer solar atmosphere: the chromosphere, the transi1on region, and the corona. •  The chromosphere is the source of most of the ultraviolet radia1on that impacts the upper atmosphere of the Earth. •  The chromosphere extends upward from the photosphere for ∼2000 km. The gas density drops to ∼10−4 that of the photosphere. Although it is almost 4 1mes thicker, this layer has much weaker absorp1on and thus lower op1cal depth over most of the visible spectrum, so that it appears only ∼10−4 as intense as the photosphere. •  The temperature varies only slightly with height, from a minimum of ∼4400 to 25,000 K, over most of the chromosphere. •  Beginning ∼2000 km above the photosphere, the density decreases rapidly by ∼3 orders of magnitude (10-­‐3 1mes less) and the temperature rises sharply to ∼106 degrees. •  The chromosphere (and at very short wavelengths, the lower corona) is revealed also in the spectral lines of the far UV, where the opacity is also high. •  At wavelengths below about 1600 Å (160 nm), the solar spectrum changes from an absorp1on to an emission spectrum, indica1ng the absence of cooler radia1on layers above the emiJng region. Figure displays the spectrum of the Sun in the so-­‐called “rocket ultraviolet” because the Earth’s atmosphere will not pass radia1on at wavelengths shorter than ∼3200 Å. The echelle spectrograph used to obtain this spectrum was carried above on an Aerobee rocket to an al1tude of ∼100 km above the Earth’s surface. •  The chromosphere is far from quiet. Every few minutes, small solar storms erupt, expelling jets of hot maUer known as spicules into the Sun’s upper atmosphere. •  These long, thin spikes of maUer leave the Sun’s surface at typical speeds of about 100 km/s, and they reach several thousand kilometers above the photosphere. •  Spicules are not spread evenly across the solar surface. Instead, they cover only about one percent of the total area, tending to accumulate around the edges of supergranules. The Sun’s magne1c field is also known to be somewhat stronger than average in those regions. The Corona •  This physically thick but extremely tenuous collec1on of gas, ionized par1cles, and dust has been likened to a thin flame. •  We are able to see the photosphere and even the chromosphere through it. •  The corona extends several solar radii out into the solar system, yet the intensity is ∼10−6 that of the photosphere. •  It is the extremely low density that makes this layer so transparent. •  The corona has three main components: the K-­‐corona, the F-­‐corona, and the E-­‐corona. •  The K-­‐corona is the light of the photosphere scaDered by electrons in the corona. This contribu1on is most important between 1 and 2.3R¤, where it dominates the coronal light. •  The F-­‐corona is the contribuFon of scaDered photospheric light by dust grains from regions beyond about 2.3R¤. This scaUering includes the absorp1on or Fraunhofer lines. The F-­‐corona, scaUered by much heavier dust par1cles, preserves the absorp1on lines beUer. •  The E-­‐corona is the source of emission lines from highly ionized atoms and is observed from all regions of the corona. The emission lines of the E-­‐corona arise from species of ions which can only appear in a very high temperature environment. •  The evidence for high temperatures in the corona comes from observa1ons of all three coronal components. Fig. IV.8 taken from “Physics and Chemistry of the Solar System”, John S. Lewis, p88 Appearance of the Solar Corona •  White-­‐light corona has loops (closed structures) and streamers (open structures), changing with the solar cycle. –  Loop-­‐like structures dominate at sunspot maximum, occurring over all la1tudes during maximum. –  Streamers are mostly concentrated around the solar equator at sunspot minimum. •  Smaller but brighter loops associated with sunspots: ac0ve regions -­‐ are seen in extreme ultraviolet and X-­‐ray emission. Coronal holes – reduced X-­‐ray emission with well-­‐defined boundaries – occur at sunspot minimum at the N and S poles. Lower-­‐la0tude holes also occur during the declining phase of the cycle. 14 The appearance of the X-­‐ray Sun throughout the solar cycle as seen by Yohkoh satellite. (Å) (Å) Spectra of the acFve sun – taken by SOHO 16 (Å) (Å) Spectra of the quiet sun – taken by SOHO 17