More About the Sun To better understand the process that creates sunspots we first need to learn more about the sun. The sun is by far the largest object in the solar system, containing more than Over time, the nuclear fusion reactions that fuel the sun are converting hydrogen into helium in its core, changing the ratio of the two elements.
Convection is the process by which hot gas from the center of the sun rises to the surface, and cooler gas, which comes to the surface and radiates its heat away, sinks back towards the center. George Fisher and David Dearborn answer the question, "What is a sunspot?
Think of a sunspot as a bubble of magnetic pressure, surrounded, by the gas pressure of the photosphere. The scale of the sun is hard to fathom. Figure 6. Because they are transparent to most visible radiation and emit only a small amount of light, these outer layers are difficult to observe. Until this century, the chromosphere was visible only when the photosphere was concealed by the Moon during a total solar eclipse see the chapter on Earth, Moon, and Sky.
Observations made during eclipses show that the chromosphere is about to kilometers thick, and its spectrum consists of bright emission lines, indicating that this layer is composed of hot gases emitting light at discrete wavelengths. The reddish color of the chromosphere arises from one of the strongest emission lines in the visible part of its spectrum—the bright red line caused by hydrogen, the element that, as we have already seen, dominates the composition of the Sun.
In , observations of the chromospheric spectrum revealed a yellow emission line that did not correspond to any previously known element on Earth.
It took until for helium to be discovered on our planet. Today, students are probably most familiar with it as the light gas used to inflate balloons, although it turns out to be the second-most abundant element in the universe. The temperature of the chromosphere is about 10, K. This means that the chromosphere is hotter than the photosphere, which should seem surprising.
In all the situations we are familiar with, temperatures fall as one moves away from the source of heat, and the chromosphere is farther from the center of the Sun than the photosphere is. Figure 7. Temperatures in the Solar Atmosphere: On this graph, temperature is shown increasing upward, and height above the photosphere is shown increasing to the right. Note the very rapid increase in temperature over a very short distance in the transition region between the chromosphere and the corona.
The increase in temperature does not stop with the chromosphere. Above it is a region in the solar atmosphere where the temperature changes from 10, K typical of the chromosphere to nearly a million degrees.
The hottest part of the solar atmosphere, which has a temperature of a million degrees or more, is called the corona.
Appropriately, the part of the Sun where the rapid temperature rise occurs is called the transition region. It is probably only a few tens of kilometers thick. Figure 7 summarizes how the temperature of the solar atmosphere changes from the photosphere outward. IRIS is the first space mission that is able to obtain high spatial resolution images of the different features produced over this wide temperature range and to see how they change with time and location Figure 8.
Figure 3 and the red graph in Figure 7 make the Sun seem rather like an onion, with smooth spherical shells, each one with a different temperature.
For a long time, astronomers did indeed think of the Sun this way. For example, clouds of carbon monoxide gas with temperatures colder than K have now been found at the same height above the photosphere as the much hotter gas of the chromosphere.
Figure 8. An image of a portion of the transition region of the corona, showing a filament, or ribbon-like structure made up of many individual threads. Like the chromosphere, the corona was first observed during total eclipses Figure 9. Unlike the chromosphere, the corona has been known for many centuries: it was referred to by the Roman historian Plutarch and was discussed in some detail by Kepler. Figure 9. Coronagraph: This image of the Sun was taken March 2, The smaller inner circle is where the Sun would be if it were visible in this image.
The corona extends millions of kilometers above the photosphere and emits about half as much light as the full moon. Just as bright city lights make it difficult to see faint starlight, so too does the intense light from the photosphere hide the faint light from the corona. While the best time to see the corona from Earth is during a total solar eclipse, it can be observed easily from orbiting spacecraft.
Studies of its spectrum show the corona to be very low in density. The corona thins out very rapidly at greater heights, where it corresponds to a high vacuum by Earth laboratory standards. These particles flow outward from the Sun into the solar system at a speed of about kilometers per second almost 1 million miles per hour! The solar wind exists because the gases in the corona are so hot and moving so rapidly that they cannot be held back by solar gravity.
This wind was actually discovered by its effects on the charged tails of comets; in a sense, we can see the comet tails blow in the solar breeze the way wind socks at an airport or curtains in an open window flutter on Earth. Although the solar wind material is very, very rarified i. Astronomers estimate that the Sun is losing about 10 million tons of material each year through this wind. While this amount of lost mass seems large by Earth standards, it is completely insignificant for the Sun.
Figure From where in the Sun does the solar wind emerge? In visible photographs, the solar corona appears fairly uniform and smooth. X-ray and extreme ultraviolet pictures, however, show that the corona has loops, plumes, and both bright and dark regions. Large dark regions of the corona that are relatively cool and quiet are called coronal holes Figure B total magnetic field strength; upperleft , vs. B r or our B h horizontal component of the magnetic field; lower-right.
As expected, the thermal brightness anti-correlates with the total field strength B since the latter is larger see Figures 4 and 11 in the darkest part of the sunspot: the umbra. Again, this was to be expected see Figures 7 and 11 since the inclination of the magnetic field increases towards the penumbra, which is brighter see Figures 2 and 3.
As we will discuss intensively throughout Section 3, these trivial results have important consequences for the energy transport in sunspots. Upper-left: scatter plot of the total field strength vs. Bottom-right: horizontal component of the magnetic field B r called B h in Section 2. In all these panels circles represent umbral points, whereas crosses and triangles correspond to points in the umbra-penumbra boundary and penumbral points, respectively from Mathew et al.
Note that in Section 1. This is, therefore, the origin of the twist: a deviation from a purely radial i. These two examples illustrate that the magnetic field vector is radial throughout most of the sunspot, but there are regions where significant deviations are observed. These deviations could be already seen in the arrows in Figures 5 and 6 representing the magnetic field vector in the plane of the solar surface. In addition, in these two examples the sign of the twist wherever it exists remains constant for the entire sunspot.
A better estimation of the twist in the magnetic field lines can be obtained from spectropolarimetric observations. To our knowledge, the first attempts in this direction were performed by Stepanov As demonstrated by Tiwari et al. Thus, any of these two parameters can be also employed to study the sign of the twist in the magnetic field vector. Using these parameters Pevtsov et al. Other possible explanations for the twist of the magnetic field in sunspots, in terms of the solar rotation and Coriolis force, have been offered by Peter and Fan and Gong However, a definite explanation is yet to be identified.
This might be more complicated than it seems at first glance because different twisting mechanisms might operate in different regimes and atmospheric layers. This is supported by recent spectropolarimetric observations that infer a twist in the magnetic field that can change sign from the photosphere to the chromosphere Socas-Navarro, a.
In this section, we will investigate the small-scale structure of the magnetic field. Another difference with Section 2, where only the magnetic field structure was discussed, is that in this section we will also address the velocity field since they are both intimately linked at small scales e. In this section spectropolarimetric observations at the highest spatial resolution will be employed and, instead of discussing averaged quantities, we will focus mostly in their pixel-to-pixel variations.
In the first part of this section we will address the fine structure of the umbral magnetic field, whereas the second will be devoted to the penumbral magnetic field. This division is somewhat artificial because the current paradigm points towards a clear relationship between the small-scale structure in these two different regions Rimmele, This difference leads to a somewhat different interaction between the convective motions and the magnetic field.
As we saw in Section 2. Plasma heated up to this temperature loses energy in the form of radiation. If the brightness of the umbra and penumbra is to remain constant, the energy loses due to radiation must be compensated by some other transport mechanism that will bring energy from the convection zone into the photosphere. The mechanism usually invoked is convection.
However, the strong magnetic field present in the sunspots see Figure 11 inhibits convective flows Cowling, This inhibited convection is, therefore, the reason why umbra and penumbra posses a reduced brightness compared to the granulation. How do these convective movements take place? In the umbra the answer to this question is to be found in the so-called umbral dots , whereas in the penumbra convection occurs within the penumbral filaments.
Umbral dots appear as small-scale regions of enhanced brightness within the umbral core see Figure Sizes and lifetimes of umbral dots have been extensively discussed in the literature. The current consensus points towards a large selection bias.
The red circles surround a local intensity enhancement in the umbral core: umbral dot UD , and a portion of a light bridge LB adapted from Henriques et al. This distinction is based upon the location of the umbral dots: CUDs appear mostly close to the darkest region of the umbra, whereas PUDs appear commonly at the umbral and penumbral boundary.
Although sometimes disputed see, e. The large continuum intensities, as compared to the umbral dark surroundings, immediately implies see, for example, Section 2. Old and current estimates all coincide in a temperature difference that ranges from K Grossmann-Doerth et al. The blue curves shows the stratification for the diffuse umbral background, where the red curves correspond to the vertical stratification along an umbral dot.
The strength of the magnetic field inside umbral dots has been a somewhat controversial subject, with some works finding no large differences between umbral dots and the umbral background Lites et al. Here the difference can be such that the magnetic field inside the umbral dot is only a few hundred Gauss see Figure This will be a recurrent topic in future sections Sections 3.
The smaller field strengths inside umbral dots leads to an enhanced gas pressure as compared to the surrounding umbra. This is consistent with the larger temperatures found inside UDs. This value is similar to the height difference for the continuum level between penumbral spines and intraspines see Figure 9. If the continuum level is actually formed higher in the geometrical height scale inside UDs than in the umbra, this means that the differences, if measured at the same geometrical height, would be much larger than the numbers previously cited.
This effect applies indeed, not only to umbral dots, but also in any other structure in the solar photosphere that is elevated with respect to its surroundings. Note also that the Wilson depression between the umbra and umbral dots can be inferred from purely geometrical considerations of sunspot observations close to the limb Lites et al.
As explained in Sect 2. As mentioned in Section 3. The main candidate for this are the umbral dots. This was motivated by the fact that umbral dots show enhanced brightness with respect to the umbral background and, therefore, must be heated more efficiently.
As it occurs in the case of penumbral filaments see Section 3. For instance, while upflows ranging from 0. However, in the past few years there have been a few positive detections of downflows at the edges of umbral dots Bharti et al. The latter work presents evidence that supports the numerical simulations of umbral convection in great detail, with umbral dots that show upflows along their central dark lane and strong downflows at the footpoints of the dark lanes see Figure 3 in Ortiz et al.
This agreement is evident if we compare the observations from Ortiz et al. Results from spectropolarimetric observations. Left panel: continuum intensity map inside the umbra of a sunspot. The circles denote the location of several umbral dots see as intensity enhancements; see also The largest circle encircles two large umbral dots that show prominent central dark lanes. Right panel: map of the line-of-sight velocity in deep layers.
This map shows an upflow blueshift along the central dark lane and downflows redshift at the footpoints of the dark lane from Ortiz et al. Results from 3D MHD simulations. Left panel: continuum intensity in the umbra of a sunspot. Right panel: map of the line-of-sight velocity. This panel shows upflows blueshift along the central dark lane of umbral dots. The lower magnetic field inside umbral dots mentioned in Section 3. In the sunspot umbra, convective motions push the magnetic field lines towards the boundary of the convective cell, thereby creating a region where the vertical component of the magnetic field vector is strongly reduced.
Since the ambient magnetic field is vertical, this automatically yields a very small field inside the umbral dot. At the top of the convective cell the magnetic field forms a cusp or canopy, preventing the material from continuing to flow upwards. Besides umbral dots, the most striking manifestation of convection in the umbra appears in the form of light bridges. These are elongated bright features that often split the umbra in two sections connecting two different sides of the penumbra see Figure Light bridges and umbral dots share many similarities.
For instance, both feature a central dark lane and bright edges. Indeed, light bridges can be considered as an extreme form of elongated umbral dots. Recent observations at much better spatial resolution have also been able to establish a clear connection between upflows and the central dark lane in light bridges, as well as downflows and the bright edges of the light bridge Hirzberger et al.
In addition, as it also occurs with umbral dots see Section 3. The nature of the central dark lane in light bridges is the same as in umbral dots Section 3.
The presence of several convective features in the umbra of sunspots immediately posses the question of whether the convective upflows and downflows in umbral dots and light bridges extend deep into the solar interior or, on the contrary, are only a surface effect. Two distinct theoretical models are usually cited to showcase these two possibilities: the cluster model Parker, and the monolithic model Gokhale and Zwaan, ; Meyer et al.
However, in the cluster model convective plumes reach very deep into the solar interior, connecting with field-free convection zone below the sunspot. In the latter model, what appears as a single flux tube in the photosphere splits into many smaller flux tubes deeper down, leaving intrusions of field-free plasma in between the smaller tubes.
Inside this intrusions is where the convection takes place. Currently, the only observational tool at our disposal that can allow us to infer the subsurface structure of sunspots is local helioseismology Gizon and Birch, ; Moradi et al. Although this technique is still under development, it will hopefully shed some light on this subject in the near future.
An alternative way of studying the subsurface structure of sunspots is by means of numerical simulations of solar magneto-convection. Furthermore, they also transport sufficient amounts of energy as to account for the observed umbral brightness see Sections 2.
At first glance these simulations seem to lend support to the monolithic sunspot model. New simulations with deeper domains have been presented by Rempel and Cheung et al. In these new simulations, umbral dots present a very similar topology as with shallower boxes. However, light bridges appear to be rooted very deep, with convective plumes that reach more than 2 Mm into the Sun see, for example, Figure 12 in Cheung et al.
Further work is, therefore, needed since the current simulations are not sufficient to completely rule out the cluster model. The filamentary structure of sunspot penumbra was recognised early in the 19th century in visual observations see review by Thomas and Weiss, The progress of observational techniques to attain higher spatial resolution revealed that the sunspot penumbra consists of radially elongated filaments with a width of 0.
Resolving the structure of the magnetic field with such high resolution is much more difficult because polarimetric measurements require multiple images taken in different polarization states, and a longer exposure time in a narrow wavelength band to isolate the Zeeman signal in a spectral line. For this reason, until recently many of the investigations of the magnetic field in the penumbra have reported contradictory results.
Wiehr and Stellmacher , however, found no general relationship between brightness and the strength of the magnetic field. Advancement of large solar telescopes at locations with a good seeing conditions made it possible to better resolve the penumbral filamentary structure in spectroscopic and polarimetric observations, and a number of papers on the small-scale magnetic field structures in sunspot penumbra were published in early s.
Lites et al. Schmidt et al. Title et al. In addition, they found no clear evidence for a spatial correlation between spines and brightness. The correlation between the magnetic field strength and field inclination i. With a highly resolved spectrum in the Fe I Better defined polarization maps of spines were taken with the Swedish 1-m Solar Telescope employing adaptive optics Langhans et al.
High quality vector magnetograms with a high spatial resolution are now routinely obtained by the spectropolarimeter SP on-board Hinode. The field inclination was derived by a Milne-Eddington inversion Section 1. It is obvious in the inclination map that the penumbrae consists of radial channels that have alternative larger and smaller field inclination. A close comparison with the continuum image shows that more horizontal field channels in panel b also called intraspines tend to be bright filaments in inner penumbra but to be dark filaments in outer penumbra.
It is confirmed that spines have stronger field than intraspines , as well as a positive field divergence. Also noticeable is the presence of a number of patches that have opposite polarity to the sunspot around the outer border of the penumbra see also Figures 4 and 7. All parameters were obtained from a Milne-Eddington inversion of the recorded Stokes spectra. The green arrow in panel a indicates the direction of the center of the solar disk. The yellow box surrounds the sunspot region displayed in Figure Thus, the penumbral magnetic field consists of two major components: spines where the magnetic field is stronger and more vertical with respect to the direction perpendicular to the solar surface, and intraspines where the magnetic field is weaker and more horizontal.
Whereas the magnetic field of the spines possibly connect with regions far from the sunspot to form coronal loops over the active region, the magnetic field in the intraspines turns back into the photosphere at the outer border of the sunspot or extend over the photosphere to form a canopy Solanki et al. The filamentary structure of the penumbra persists even after averaging a time series of continuum images over 2—4.
This suggests that the two magnetic field components are more or less exclusive to each other Thomas and Weiss, ; Weiss, except for a possible interaction through reconnection at the interface between them in the photosphere Katsukawa et al.
Such structure of the penumbral magnetic field, i. The fact that the magnetic field is weakened in the intraspines , as compared with the spines , can also be employed to deduce through total pressure balance considerations as we already did in the case of umbral dots and light bridges; see Section 3. To account for the filamentary structure of penumbra with the uncombed magnetic fields, some distinguished models, that are under a hot discussion nowadays, were proposed.
One of these models, the embedded flux tube model is an empirical model proposed by Solanki and Montavon , in which nearly horizontal magnetic flux tubes forming the intraspines are embedded in more vertical background magnetic fields spines in the penumbra Figure 25 , left panel.
The downward pumping mechanism Thomas et al. In this scenario, submergence of the outer part of flux tubes occurs as a result of the downward pumping by the granular convection outside the sunspots, and such magnetic fields form the low-laying horizontal flux tubes. Another idea to account for the penumbral filaments is the field-free gap model initially proposed by Choudhuri and later refined by Spruit and Scharmer and Scharmer and Spruit Here, the penumbral bright filaments are regarded as manifestations of the protrusion of non-magnetized, convecting hot gas into the background oblique magnetic fields of the penumbra.
This immediately yields a region, above the penumbral filaments, where the magnetic field is almost horizontal i. Models for explaining the uncombed penumbral structure. Upper-left: embedded flux tube model from Solanki and Montavon, , reproduced by permission of the ESO ; lower-left: rising flux tube model from Schlichenmaier et al. All the aforementioned models attempt to explain, with different degrees of success, the configuration of the magnetic field in the penumbra.
However, the appearance of a penumbra is always associated with a distinctive gas flow, i. In the next section we will address this issue. The Evershed flow was discovered in by John Evershed at the Kodaikanal Observatory in India as red and blue wavelength shifts in the spectra of absorption lines in the limb-side and disk-centerside of the penumbra, respectively. This feature can be explained by a nearly horizontal outflow in the photosphere of the penumbra Evershed, An outstanding puzzle about the Evershed flow lies in the relation between the velocity vector and the magnetic field vector in the penumbra.
It is highly plausible that there is a close relationship between the Evershed flow and the filamentary structure of the penumbra. Indeed, it was recognized in the s that the flow is not spatially uniform but concentrated in narrow channels in penumbra; e.
Two models were proposed to account for the nature of penumbral filaments and the Evershed flow before One is the elevated dark filament model in which the penumbral dark regions are regarded as elevated fibrils with nearly horizontal magnetic field overlaying the normal photosphere and carry the Evershed flow in them Moore, ; Cram and Thomas, ; Thomas, ; Ichimoto, The other is the rolling convection model in which penumbral filaments are regarded as convective elements radially elongated by a nearly horizontal magnetic field in penumbra and where the Evershed flow is confined in dark lanes that are analogous to the intergranular dark lanes Danielson, b ; Galloway, The long-lasting enigma on the Evershed flow was finally solved by the discovery of the interlocking comb structure of the penumbral magnetic field Section 3.
Under this scenario, the Evershed flow is confined in nearly horizontal magnetic field channels in penumbra i. Both components, when averaged together, make the spatially averaged magnetic field far from completely horizontal Figures 11 and 26 ; see also Title et al. The relationship between the Evershed flow and the horizontal magnetic field in the penumbra has been highlighted in many works in the past: Stanchfield II et al.
Geometry of magnetic field and Evershed flow in penumbra. Magnetic field lines are shown by inclined and colored cylinders, while the Evershed flow is indicated by white arrows in dark penumbral channels.
Note that the Evershed flow concentrates along the more horizontal magnetic field lines white cylinders from Title et al. Dotted curves in each panel show the scattered light fraction obtained from the inversion algorithm. Note that the velocity Evershed flow is strongest in the regions where the magnetic field is weak and horizontal intraspines , while it avoids the regions with more less inclined and stronger magnetic field spines from Borrero and Solanki, , reproduced by permission of the AAS.
In the embedded flux-tube model Solanki and Montavon, , the Evershed flow is supposed to be confined in the horizontal magnetic flux tubes embedded in more vertical background magnetic field of the penumbra. In such picture, the siphon flow mechanism was proposed as the driver of the Evershed flow Meyer and Schmidt, ; Thomas, ; Degenhardt, ; Montesinos and Thomas, : a difference in the magnetic field strength between two footpoints of a flux tube causes a difference of gas pressure, and drives the flow in a direction towards the footpoint with a higher field strength i.
Schlichenmaier et al. By performing an inversion of Stokes profiles of three infrared spectral lines at nm and using a two component penumbral model in which two different magnetic atmospheres are interlaced horizontally, Bellot Rubio et al. This picture was supported by Borrero et al. With a further elaborated analysis, Borrero et al. They argued that these results strongly support the siphon flow as the physical mechanism responsible for the Evershed flow.
Until recently the relationship between the Evershed flow and the brightness of the penumbral filaments has been somewhat controversial. Many authors Beckers, ; Title et al. Rimmele a showed that the correlation becomes better when one compares the intensity and velocity originating from the same height, and also gave a hint that the correlation is different between inner and outer penumbra. Figure 28 presents the spatial correlation between penumbral filaments and the Evershed flow.
The correlation coefficient between the Doppler shift V los and the elevation angle of magnetic field vector from the solar surface Footnote 9 as a function of the radial position in the penumbra, is displayed in the upper-right panel, whereas the correlation between the Doppler shift V los and continuum intensity is shown in the lower-right panel. In these plots, the results for limb-side penumbra are shown in red color and for disk-center-side are shown in blue color.
The abscissa in this figure spans from the umbra-penumbra boundary left to the outer border of the penumbra right. The curves along which the correlation coefficients are obtained are shown for both limb-side and disk-center-side penumbra in the left panels. In this plot, line-of-sight velocities are taken in absolute value such that there is no difference between the redshifts in the limb-side and the blueshifts in the center side that are characteristic of the Evershed flow.
In Figure 28 , we notice that the Evershed flow correlates with more horizontal magnetic fields throughout the entire penumbra, while it correlates with bright filaments in the inner penumbra but with dark filaments in the outer penumbra.
These results are consistent with the idea that penumbral filaments, which harbor a nearly horizontal magnetic field, are brighter in inner penumbra but darker in outer penumbra. The correlation between Doppler shift and intensity shows an asymmetric distribution between the disk center-side and limb-side penumbra lower right. This suggests that overposed to the Evershed flow, which is mainly horizontal, there exists a vertical component in the velocity vector in the penumbra. This vertical component will be discussed in detail in Sections 3.
Spatial correlation between penumbral filaments and the Evershed flow. Correlation coefficients between the Evershed flow and brightness lower-right , and between the Evershed flow and the elevation angle of magnetic field from the solar surface upper-right are shown as a function of radial distance from the sunspot center. Arc-segments along which the correlation is calculated are shown in the left panel. The direction to the center of the solar disk is shown by an arrow in the left panel.
Red lines show the results for the limb-side penumbra, whereas blue corresponds to the center-side penumbra. Some results from the latter two works are presented in Figures 29 and Depth structure of penumbra derived from Stokes inversions of spectro-polarimetric data.
Showns are vertical cuts across the penumbral filaments. The horizontal axis is the azimuthal direction around the penumbra and, therefore, it is perpendicular to the radial penumbral filaments. Upper-left panel : line-of-sight velocity V los. Upper-right : total magnetic field strength B. This plot demonstrates that the strong and vertical magnetic field of the spines extends above the intraspines indicated by the index i , where the Evershed flow is located where the magnetic field is rather horizontal and weak.
It also shows that the azimuth of the magnetic field changes sign above the intraspines, indicating that the magnetic field of the spines wraps around the intraspines. The arrows in this figure show the direction of the magnetic field in the plane perpendicular to the axis of the penumbral filaments from Borrero et al.
One important issue that needs to be addressed to understand the origin of the penumbra is how the energy transport takes place.
In search for these convective motions we have to examine the predictions that the different models see Section 3. The hot rising flux-tube model see Section 3. Provided that the flux tubes are evenly distributed, this would yield a preference for upflows in the inner penumbra, whereas downflows would dominate at large radial distances. Because of these features, we will refer to this form of convection as radial convection , with convective flows occurring along the penumbral filament see upper panel in Figure This implies that there should be a significant magnetic flux and mass flux returning to and emerging from the photosphere in penumbra.
Possible patterns of convection present in the sunspot penumbra. The upper panel corresponds to a pattern of radial convection, where upflows are presented at the inner footpoints of the penumbral filaments and downflows at the outer footpoints. This pattern is predicted by the embedded flux-tube model and the hot rising flux-tube model.
This is the flow pattern predicted by the field-free gap model. This strongly resembles to the convective motions discussed in Sections 3. The field-free gap model does not predict any particular radial preference for upflows and downflows in the penumbra.
Because of this feature we will refer to this type of convection as azimuthal convection or overturning convection see lower panel in Figure Note that the field-free gap model does not readily offer an explanation for the Evershed flow. This is an important point that will be addressed in Section 3. In order to distinguish between the different proposed models that attempt to explain the heating mechanism in the penumbra Section 3.
Since its discovery, the Evershed flow has been recognized as a horizontal motion of the photospheric gas. However, as already mentioned in Section 3.
A clear evidence for the vertical motions appeared only after the s when high spatial resolution became available in spectroscopic observations. Since then, a number of observations have been reported regarding the vertical component of the flow in the penumbra. On the one hand, upflows in the penumbra have been reported by Johannesson , Schlichenmaier and Schmidt , , and Bellot Rubio et al. On the other hand, downflows have been observed in and around the outer edge of penumbra by, among others, Rimmele b , Westendorp Plaza et al.
Both down- and upflows have been simultaneously observed in the penumbra by Schmidt and Schlichenmaier , Schlichenmaier and Schmidt , Westendorp Plaza et al. Figure 32 highlights some observations that clearly show the downflow patches around a sunspot with an opposite polarity Westendorp Plaza et al.
These last features correspond to the solar called bright penumbral grains and they are related to the peripheral umbral dots Sobotka et al. Selected observations of vertical motions in sunspots. Left panels : Discovery of downflows around the outer border of a sunspot. The sunspot is located near the center of the solar disk. Top is the continuum image and bottom is the magnetic field inclination overlaid with velocity contours.
Right panels : Close-up of the inner part of a limb-side penumbra. Each Evershed flow channel white filaments in the Dopplergram is associated with a bright grain and upflow dark point in the Dopplergram from Rimmele and Marino, , reproduced by permission of the AAS. According to the picture drawn by the present observations, most material carried by the Evershed flow is, thus, supposed to flow back into the photosphere at the downflow patches Westendorp Plaza et al.
The configuration of the magnetic field is affected by the aforementioned vertical flows. In fact, some of the magnetic field lines plunge back into the deep photosphere at the outer edge of the sunspot and its surroundings see Figures 4 and 7. The relationship between vertical motions and the magnetic field vector in the penumbra is clearly demonstrated by spectropolarimetric data of a sunspot near disk center. If there are no mass motions in the sunspot, Stokes V maps in the blue and red wings are expected to be identical since the Zeeman effect produces anti-symmetric Stokes V profiles around the line center.
As is confirmed by the Dopplergram in the line-wing of Stokes I bottom-left panel in Figure 33 , the former features are associated with upward motions while the later correspond to strong downflows. The presence of very fast downflows in the mid and outer regions of the penumbra has been reported previously by del Toro Iniesta et al.
Bottom panels : line-of-sight velocity Doppler velocity measured in the wings left and on the core right of the spectral line. It is obvious in panel b that the upflow and downflow patches are aligned with nearly horizontal field channels filaments with light appearance in the inclination map that carry the Evershed flow, and that small-scale upflows are preferentially located near the inner penumbra, while downflows dominate at the outer ends of the horizontal field channels see also Figure 1 in Ichimoto, and Figure 5 in Franz and Schlichenmaier, Thus, the upflow and downflow patches seen here can be regarded as the sources and sinks of the elementary Evershed flow embedded in deep penumbral photosphere.
When all the aforementioned results and observations for the velocity and the magnetic field vector are put together, the picture of the penumbra that emerges is that of Figure Overlaid are contours for upflow regions with 0. These results seem to support the idea that convective motions occur in a radial pattern along the penumbral filaments radial convection ; see upper panel in Figure In principle, this lends a strong support to the hot rising flux-tube model see Section 3.
This would imply that the energy carried by the upflows is not sufficient to heat the penumbra according to the argument of Schlichenmaier and Solanki Indeed, the search for an azimuthal convective pattern, i.
Unfortunately, results based on spectroscopic measurements have been contradictory so far. Whereas some works Rimmele, ; Zakharov et al. One of the reasons is the lack of sufficient spatial resolution to resolve the velocity fields inside the penumbral filaments. Regardless of which form of convection takes place, it is very clear that this is indeed the mechanism that is responsible for the energy transport in the penumbra. This is emphasized by the very close relationship existing between upflows and bright grains in the penumbra as seen in Figure The blue contours show blueshifts equal or larger than 0.
The fact that upflows blue contours correlate so well with bright penumbral regions, strongly suggest that the vertical component of the Evershed flow supplies the heat to maintain the penumbral brightness even though a quantitative evaluation of the heat flux is not available Ichimoto et al. Puschmann et al. Figure 36 shows typical Stokes profiles in upflow panel a and downflow panel b regions in the penumbra, respectively. It is noticeable that the upflow region shows a blue hump in Stokes V with the same polarity of the main lobe in the blue wing, while the downflow region shows a strong third lobe with opposite polarity in the far red wing of Stokes V.
The solid curves show the best-fit profiles produced by a Milne-Eddington ME inversion algorithm. A ME-inversion assumes that the physical parameters are constant with optical depth see Section 1.
For instance, Sainz Dalda and Bellot Rubio found small-scale, radially elongated, bipolar magnetic structures in the mid-penumbra aligned with intraspines. They move radially outward and were interpreted by these authors as manifestations of the sea-serpent field lines that harbor the Evershed flow Schlichenmaier, and, eventually, leave the spot to form moving magnetic features.
Solid curves show results of a Milne-Eddington fitting algorithm see Section 1. Improvements in the spatial resolution in ground-based optical observations revealed further details about the rich variety of fine-scale structures in the penumbra.
Scharmer et al.
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