During the total solar eclipse of July 22, 2009, we carried out a white-light observation in Anji, Zhejiang, China. The aim wasto observe the polar plumes (PPs) with high spatial and temporal resolutions in white-light. With the observational data, weinvestigate the properties and evolution of the PPs and compare them with those of the low-latitude plumes (LPs). We find thatboth the PPs and the LPs have comparable lengths and widths, and the mean length and width are 300 Mm and 16 Mm, re-spectively. The average inclination angle (13 degree) of the PPs is smaller than that (32 degree) of the LPs. Generally, theplumes which are closer to the coronal hole center are more vertical. We trace the PPs and the LPs in the sequence of imagesand find that none of them disappears and no new one is created. Additionally, neither plasma outflow nor transverse oscilla-tion is observed. These imply that the evolution process of plumes is much longer than the timescale of eclipse.
Using the multi-wavelength data from the Atmospheric Imaging Assembly （AIA） onboard the Solar Dynamics Observatory （SDO） spacecraft, we study a jet occurring in a coronal hole near the northern pole of the Sun. The jet presented distinct upward helical motion during ejection. By tracking six identified moving features （MFs） in the jet, we found that the plasma moved at an approximately constant speed along the jet＇s axis. Meanwhile, the MFs made a circular motion in the plane transverse to the axis. Inferred from linear and trigonometric fittings to the axial and transverse heights of the six tracks, the mean values of the axial velocities, transverse velocities, angular speeds, rotation periods, and rotation radii of the jet are 114 km s-1, 136 km s-1, 0.81° s-1, 452 s and 9.8 × 10^3 km respectively. As the MFs rose, the jet width at the corresponding height increased. For the first time, we derived the height variation of the longitudinal magnetic field strength in the jet from the assumption of magnetic flux conservation. Our results indicate that at heights of 1 × 10^4 -7 × 10^4 km from the base of the jet, the flux density in the jet decreases from about 15 to 3 G as a function of B = 0.5（R/R） - 1）-0.84 （G）. A comparison was made with other results in previous studies.
With the observations of the Solar-Terrestrial Relations Observatory （STEREO） and the Solar Dynamics Observatory （SDO）, we analyze in detail the kine- matics of global coronal waves together with their intensity amplitudes （so-called ＂perturbation profiles＂）. We use a semi-automatic method to investigate the pertur- bation profiles of coronal waves. The location and amplitude of the coronal waves are calculated over a 30~ sector on the sphere, where the wave signal is strongest. The position with the strongest perturbation at each time is considered as the location of the wave front. In all four events, the wave velocities vary with time for most of their lifetime, up to 15 rain, while in the event observed by the Atmospheric Imaging Assembly there is at, additional early phase with a much higher velocity. The velocity varies greatly between different waves from 216 to 440 km s-1. The velocity of the two waves initially increases, subsequently decreases, and then increases again. Two other waves show a deceleration followed by an acceleration. Three categories of am- plitude evolution of global coronal waves are found for the four events. The first is that the amplitude only shows a decrease. The second is that the amplitude initially increases and then decreases, and the third is that the amplitude shows an orderly in- crease, a decrease, an increase again and then a decrease. All the extreme ultraviolet waves show a decrease in amplitude while propagating farther away, probably because the driver of the global coronal wave （coronal mass ejection） is moving farther away from the solar surface.
Solar activity,in which there is an explosive release of magnetic energy in the solar atmosphere,is one of the most spectacular manifestations of space plasma activity.Non-potential,free magnetic energy is accumulated by the interaction between plasma and magnetic fields in the photosphere and sub-photospheric layer.In the photosphere and lower solar atmosphere,plasma is only partially ionized and there is three-species(or three-fluid)plasma.Cowling conductivity should therefore be considered in explaining solar observations.On the other hand,the explosive release of magnetic energy is believed to be caused by magnetic reconnection in the corona,where plasma is fully ionized and essentially collisionless.However,for collisionless plasma,we seem to have no precise quantitative description or formulation of its conductivity.To physically understand magnetic reconnection in collisionless plasma is an important task in making further progress in solar activity studies.This article discusses the plasma perspective of solar activity studies.The discussion begins with solar observations,and then,an analysis of the plasma problems that we face and need to better understand.
Magnetic non-potentiality is important for understanding flares and other solar activities in active regions （ARs）. Five non-potential parameters, i.e. electric current, current helicity, source field, photospheric free energy, and angular shear, are calculated to quantify the non-potentiality of NOAA AR 11158. Benefitting from the high spatial resolution, high cadence and continuous temporal coverage of vector mag- netograms from the Helioseismic and Magnetic Imager onboard the Solar Dynamics Observatory, both the long-term evolution of the AR and the rapid change during flares are studied. We confirm that, compared with the magnetic flux, the magnetic non-potentiality has a closer connection with the flare, and the emerging flux regions are important for understanding the magnetic non-potentiality and flares. The main re- suits are as follows. （1） The vortex in the source field directly displays the deflection of the horizontal magnetic field. The deflection corresponds to the fast rotating sunspot with a time delay, which suggests that the sunspot rotation leads to an increase in the non-potentiality. （2） Two areas that have evident changes in the azimuth of the vector magnetic field are found near the magnetic polarity inversion line. The change rates of the azimuth are about 1.3° h-1 and 3.6° h-1, respectively. （3） Rapid and prominent increases are found in the variation of helicity during four flares in the regions where their initial brightening occurs. The recovery of the increases takes 3-4 h for the two biggest flares （X2.2 and M6.6）, but only takes about 2 h for the two other smaller flares （M2.2 and M1.6）.
