Upflows in the upper solar atmosphere
Hui Tian, Louise Harra, Deborah Baker, David H. Brooks, Lidong Xia
SSolar PhysicsDOI: 10.1007/ ••••• - ••• - ••• - •••• - • Upflows in the Upper Solar Atmosphere
Hui Tian · Louise Harra · Deborah Baker · David H. Brooks · Lidong Xia © Springer ••••
Abstract
Spectroscopic observations at extreme- and far-ultraviolet wavelengthshave revealed systematic upflows in the solar transition region and corona. Theseupflows are best seen in the network structures of the quiet Sun and coronal holes,boundaries of active regions, and dimming regions associated with coronal massejections. They have been intensively studied in the past two decades becausethey are likely to be closely related to the formation of the solar wind and heatingof the upper solar atmosphere. We present an overview of the characteristicsof these upflows, introduce their possible formation mechanisms, and discusstheir potential roles in the mass and energy transport in the solar atmosphere.Although past investigations have greatly improved our understanding of theseupflows, they have left us with several outstanding questions and unresolvedissues that should be addressed in the future. New observations from the
SolarOrbiter mission, the
Daniel K. Inouye Solar Telescope and the
Parker SolarProbe will likely provide critical information to advance our understanding ofthe generation, propagation, and energization of these upflows. (cid:66)
H. [email protected] School of Earth and Space Sciences, Peking University, Beijing 100871, China Key Laboratory of Solar Activity, National Astronomical Observatories, ChineseAcademy of Sciences, Beijing 100012, China PMOD/WRC, Dorfstrasse 33, 7260 Davos Dorf, Switzerland ETH-Z¨urich, H¨onggerberg Campus, Z¨urich, Switzerland Mullard Space Science Laboratory, University College London, Holmbury, St. Mary,Dorking, Surrey, KT22 9XF, UK College of Science, George Mason University, 4400 University Drive, Fairfax, VA22030, USA Shandong Provincial Key Laboratory of Optical Astronomy and Solar-TerrestrialEnvironment, Institute of Space Sciences, Shandong University, Weihai, 264209Shandong, China
SOLA: upflow.tex; 9 February 2021; 2:31; p. 1 . Tian et al.
Keywords:
Active Regions, Velocity Field; Coronal Holes; Coronal Mass Ejec-tions, Low Coronal Signatures; Heating, Coronal; Spectral Line, Broadening
1. Introduction
The upper solar atmosphere consists of the corona and transition region (TR),spanning a temperature range from about 3 × to several million Kelvin. Insuch a hot environment, atoms are often highly ionized and produce hundredsof strong emission lines mainly at extreme-ultraviolet (EUV) and far-ultraviolet(FUV) wavelengths ( ≈
100 – 1700 ˚A). In the past quarter century, dedicated ob-servations with several EUV/FUV spectrographs, particularly the
EUV ImagingSpectrometer (EIS: Culhane et al., 2007) onboard
Hinode (Kosugi et al., 2007),the
Interface Region Imaging Spectrograph (IRIS: De Pontieu et al., 2014), the
Solar Ultraviolet Measurements of Emitted Radiation (SUMER: Wilhelm et al.,1995; Lemaire et al., 1997), and the
Coronal Diagnostic Spectrometer (CDS:Harrison et al., 1995) onboard the
Solar and Heliospheric Observatory (SOHO:Domingo, Fleck, and Poland, 1995), have greatly improved our understandingof various types of dynamic activity in the upper solar atmosphere.Signatures of outflows or upflows, i.e. blue shifts of spectral lines, have beenfrequently reported from spectroscopic observations with these facilities. In thisreview, we focus on the systematic and pervasive upflows observed primarily withthe aforementioned spectrographs. Sporadic upflows/outflows such as coronalmass ejections and coronal jets are not discussed in this review. These systematicupflows have been one focus of investigation in the past two decades, mainly be-cause they are likely to be related to the formation of the solar wind. In addition,these upflows are expected to contribute to the mass and energy transport inthe solar atmosphere, and thus may play an important role in coronal heating.Here we present a review on both observational and theoretical investigationsof these upflows. We first discuss upflow signatures in the quiet Sun and coronalholes in Section 2, then provide a detailed introduction to the coronal-upflowphenomenon at active-region (AR) boundaries in Section 3. Section 4 describescharacteristics of the upflows identified from regions of coronal dimmings inducedby coronal mass ejections (CMEs). Finally, we summarize the major findingsabout these upflows and briefly discuss future perspectives in Section 5.
2. Upflows from the Quiet Sun and Coronal Holes
It has been known since the 1970s that emission lines formed in the TR, such asC iv ≈ K under ionization equilibrium, areredshifted by a few km s − on average in the quiet Sun (e.g. Doschek, Feldman,and Bohlin, 1976). The red shift is obviously larger in network regions comparedto internetwork regions (Figure 1). With the capability of observing hundredsof strong emission lines formed in a wide range of temperatures, the SUMER SOLA: upflow.tex; 9 February 2021; 2:31; p. 2 pflows
Figure 1.
Intensity and Doppler shift images of C iv viii
770 ˚A. Blue and redcolors in the Dopplergrams indicate blue shifts and red shifts, respectively. These images corre-spond to a field of view (FOV) of 540 (cid:48)(cid:48) × (cid:48)(cid:48) spectrometer allowed detailed investigations of the Doppler shift as a functionof the line-formation temperature. SUMER observations revealed a clear depen-dence of the Doppler shift on temperature, i.e. the average red shift increaseswith temperature and peaks around 2 × K (or log T = 5 .
3) (e.g. Peter andJudge, 1999; Stucki et al., 2000; Dadashi, Teriaca, and Solanki, 2011). As thetemperature continues to increase, the average red shift decreases and turns intoa blue shift above ≈ × K. This trend can be clearly seen from Figure 2.The change from downflow (i.e. red shifts of spectral lines) dominance to up-flow dominance was unambiguously established from observations of the strongNe viii
770 ˚A line, which is formed at ≈ × K in the upper TR or lowercorona. Under the assumption of zero average Doppler shift above the solar limb,Dammasch et al. (1999) determined a rest wavelength of 770.428 ± SOLA: upflow.tex; 9 February 2021; 2:31; p. 3 . Tian et al.
Figure 2.
Average Doppler shifts of several spectral lines with different formation tempera-tures (Xia, Marsch, and Wilhelm, 2004). (a) A coronal hole region. (b) A quiet-Sun region. (c)Difference of the Doppler shifts between the coronal hole and quiet-Sun regions. Reproducedwith permission from
Astronomy & Astrophysics , © ESO. this line. The same rest wavelength was also found independently by Peter andJudge (1999). Using this rest wavelength, the Ne viii
770 ˚A line was found tobe blueshifted by ≈ − on average (e.g. Peter and Judge, 1999; Teriaca,Banerjee, and Doyle, 1999; Xia, Marsch, and Wilhelm, 2004). Dopplergrams ofNe viii
770 ˚A obtained through raster scans revealed very prominent blue shiftsat junctions of multiple adjacent network structures (e.g. Hassler et al., 1999;Dammasch et al., 1999; Aiouaz, 2008; Tian et al., 2008). These localized blueshifts could reach 5 – 10 km s − (Figure 1). They generally correspond to thelegs of magnetic loops reconstructed through force-free magnetic-field extrap- SOLA: upflow.tex; 9 February 2021; 2:31; p. 4 pflows olations, indicating mass supply to the corona along flux tubes rooted in thechromospheric network (Tian et al., 2009). Using the Doppler shift of Ne viii asa proxy for the plasma bulk flow (i.e. the proton flow), these authors also foundan anti-correlation between the flux tube expansion factor and mass flux.
A similar dependence of Doppler shift on formation temperature has also beenfound in coronal holes. From Figure 1 we can see that the C iv iv Dopplergrams in the quiet Sun and coronal holes, i.e.there are more pixels with blue shifts in a coronal hole than in an equally-sizedquiet-Sun region (Dammasch et al., 1999). Figure 1 shows that Ne viii
770 ˚A isblueshifted almost everywhere in a polar coronal hole, which is also distinctlydifferent from the localized blue shifts of Ne viii in the quiet Sun. The preva-lence of blue shifts has been found in both polar and equatorial coronal holes(Hassler et al., 1999; Wilhelm et al., 2000; Xia, Marsch, and Curdt, 2003; Tuet al., 2005a; Aiouaz, Peter, and Lemaire, 2005). These prevalent blue shifts havebeen observed in not only the Ne viii
770 ˚A line, but also in higher-temperaturecoronal lines such as Fe x xii T = 4 . T = 6 . T ≈ . ≈ × K in coronal holes are generally believed to be signatures of the nascentfast solar wind (e.g. Hassler et al., 1999; Wilhelm et al., 2000; Xia, Marsch,and Curdt, 2003). The three-dimensional (3D) coronal magnetic field could bereconstructed from a measured photospheric magnetogram through a potentialor force-free field extrapolation, and it can thus be used to correlate with theplasma flow pattern (Wiegelmann, Xia, and Marsch, 2005). Using a similarmethod, Tu et al. (2005a) identified many funnel-like magnetic-flux tubes rootedin the chromospheric network of a coronal-hole region, and found that patchesof large Ne viii blue shift are closely associated with these funnels, suggestingthat the nascent fast solar-wind flows outward along these expanding magneticfunnels (Figure 3). They calculated the correlation coefficient between maps ofthe Ne viii blue shift and magnetic-field inclination at each height, and theyfound a maximum correlation at a height of ≈
20 Mm above the photosphere.This height could be regarded as a rough estimate of the real formation heightof Ne viii
770 ˚A, because the motion of the emitting ions is expected to be largelyguided and controlled by the shape of the expanding funnels at the line-emissionheight in the low- β environment. SOLA: upflow.tex; 9 February 2021; 2:31; p. 5 . Tian et al.
Figure 3.
Solar wind origin from funnel-like magnetic structures (Tu et al., 2005b). Lowerleft: a full-disk coronal image. Top: reconstructed 3D magnetic-field structures in the whiterectangular region. The black and purple lines represent closed- and open-field lines, respec-tively. A photospheric magnetogram (radial component, blue and red colors indicate differentpolarities) is shown at the bottom. The higher plane shows the magnetogram at 20 Mm, withthe hatched region indicating largest blue shifts ( (cid:62) − ) of Ne viii
770 ˚A. Lower right:zoomed-in view of the region indicated by the red rectangle.
It is worth mentioning that apparent jet-like features at coronal temperatureshave been frequently reported in coronal holes (e.g. Cirtain et al., 2007; Ni et al.,2020; Shen, 2021). Many of these jets are associated with coronal bright points. Inaddition, Young (2015) found a new type of jet that is visible in the spectroscopicdata but shows no clear signature in the imaging data. They are rooted in brightpoints and have speeds reaching ≈
100 km s − . These sporadic jets are possiblynot the dominant cause of the net blue shifts. The observed temperature dependence of the Doppler shift is still not well un-derstood. Tu et al. (2005a) and He, Tu, and Marsch (2008) interpreted the blueand red shifts in coronal holes as the bidirectional flows generated by magneticreconnection between the open-field lines in coronal funnels and adjacent low-lying loops. A similar continuous reconnection between legs of large coronal loops(strong network field) and low-lying loops (weak field in internetwork regions)might also exist in the quiet Sun, resulting in bidirectional flows that accountfor the blue shifts of hotter lines and red shifts of cooler lines (Aiouaz, 2008).An alternative interpretation, which has received much more attention inthe community, considers the mass-circulation process in the solar atmosphere.
SOLA: upflow.tex; 9 February 2021; 2:31; p. 6 pflows
Figure 4.
Images of N v − (from line centroid) spectral range at the blue and red wings, respectively. Thecontours outline regions of strong intensity. The FOV of each image is 90 (cid:48)(cid:48) × (cid:48)(cid:48) . Reproducedby permission of the AAS. Although spectral lines with different formation temperatures show differentsigns of net Doppler shifts, a careful examination of the TR spectral-line profilesobserved by SUMER and IRIS in network regions suggests that the spectral pro-files are often enhanced at the blue wing (Figure 4). The blue-wing enhancementor blueward asymmetry has been found to occur intermittently, suggesting thepresence of episodic weak ( ≈ ≈
50 – 120km s − ) upflows (McIntosh and De Pontieu, 2009a; De Pontieu et al., 2009; Chenet al., 2019b). SUMER observations show that the weak blueward asymmetrycan be identified from spectral lines formed in a wide range of temperatures, atleast from ≈ × K to ≈ × K (McIntosh and De Pontieu, 2009a). Theseauthors interpreted the observed emission of the TR and lower corona as amixture of emission from rapid injection of episodically heated plasma to thecorona and slow cooling of the previously heated plasma. They conjecturedthat the high-speed plasma ejections could be heating signatures of the high-speed spicules observed in the chromosphere (De Pontieu et al., 2007). Withsimultaneous imaging and spectroscopic observations of the TR, Tian et al.(2014b) discovered prevalent high-speed intermittent jets with a temperatureof at least ≈ K, and they demonstrated that the blueward asymmetry ofthe Si iv SOLA: upflow.tex; 9 February 2021; 2:31; p. 7 . Tian et al. relative contributions of the three components at different temperatures mightbe responsible for the observed temperature dependence of Doppler shift. Themore or less steady behavior of the observed Doppler shifts suggests that sucha mass-cycling process should occur continuously.Numerical simulations have also been performed to understand the temper-ature dependence of Doppler shift. Gudiksen and Nordlund (2005) developed a3D magnetohydrodynamic (MHD) model of the solar corona, where the coronalheating is due to the magnetic braiding induced by field line footpoint motionsin the photosphere. Peter, Gudiksen, and Nordlund (2004) and Peter, Gudiksen,and Nordlund (2006) took this model and synthesized spectra of several TR andcoronal lines. They found that this model can reproduce the observed temper-ature dependence of the TR red shift, implying that the well-known TR redshifts are caused by the flows induced by heating through braiding of magnetic-field lines. However, the model prodicted a net red shift for the Ne viii
770 ˚A andMg x
625 ˚A lines that are formed in the upper TR and lower corona, which isinconsistent with SUMER observations but might be due to the relatively lowheight of the upper boundary. Using 3D MHD models spanning the upper con-vection zone to 15 Mm above the photosphere, Hansteen et al. (2010) found thatrapid intermittent heating of the plasma from the upper chromosphere to coronaltemperatures naturally produces net red shifts for most TR lines and smallblue shifts of spectral lines formed above a temperature of ≈ × K, roughlyconsistent with SUMER and EIS observations. These episodic heating events areaccompanied by the generation of intermittent high-speed multi-thermal upflows,which may be related to the fast chromospheric spicules (De Pontieu et al.,2007), TR network jets (Tian et al., 2014b), and high-speed upflows inferredfrom the blueward asymmetries of TR lines observed by SUMER (McIntosh andDe Pontieu, 2009a). From high-resolution measurements of the photosphericmagnetic field, Samanta et al. (2019) demonstrated that intermittent magneticreconnection between the strong network field and the newly emerged, weak,small-scale internetwork field is likely responsible for the generation of at leastsome of these upflowing materials.
3. Upflows from Active-Region Boundaries
The dominant emission structures in the upper solar atmosphere of ARs arelarge-scale loops that outline the strong magnetic field. Prominent red shifts ofspectral lines formed in the TR are commonly observed in these coronal loops,particularly at the loop legs (e.g. Winebarger et al., 2002; Marsch, Wiegelmann,and Xia, 2004; Dammasch et al., 2008; Marsch et al., 2008; Del Zanna, 2008).Different from the TR red shifts observed in the network lanes of quiet-Sunregions and coronal holes, these red shifts are also present in coronal lines witha formation temperature as high as ≈ . K (Del Zanna, 2008). The magnitudeof these red shifts could reach ≈
30 km s − , two to three times higher than inthe quiet Sun and coronal holes. Despite these differences, the origin of thesered shifts is likely to be similar to that in the quiet Sun and coronal holes (e.g.Dammasch et al., 2008). SOLA: upflow.tex; 9 February 2021; 2:31; p. 8 pflows
At the boundaries of many ARs, the coronal emission is normally weaker thanin the AR cores. With SOHO/SUMER observations, Marsch, Wiegelmann, andXia (2004) found signatures of upflows in the Ne viii
770 ˚A line, which is formedin the upper TR or lower corona, at the boundaries of ARs. With
Hinode /EIS ob-servations, coronal spectral lines such as Fe xii xiii
Hinode
ReviewTeam et al., 2019). In this section we provide an updated and more extendedreview of past investigations on these upflows.
These blue shifts appear to be quasi-steady, i.e. they generally last for at leasta few days (e.g. Baker et al., 2009; Bryans, Young, and Doschek, 2010; Tianet al., 2012b). A single Gaussian fit (SGF) to the observed coronal line profilesoften leads to a blue shift of ≈
10 – 50 km s − (e.g. Harra et al., 2008; Marschet al., 2008; Doschek et al., 2008; Srivastava et al., 2014). These blue shifts areassociated with enhanced line broadenings (or nonthermal velocities), and thereis a positive correlation between them (e.g. Doschek et al., 2008). From magnetic-field extrapolations, these upflows appear to be guided by the legs of large-scalemagnetic loops or open magnetic-field structures at AR boundaries (e.g. Marsch,Wiegelmann, and Xia, 2004; Harra et al., 2008; Marsch et al., 2008; Baker et al.,2009). These structures expand rapidly with height, often appearing as fan-likestructures. The upflows are normally more obvious towards the footpoints of thefans, but are not always associated with them (Warren et al., 2011). They are,however, sometimes very prominent in low-emission regions at the peripheriesof AR cores or between the cores and fans (e.g. McIntosh et al., 2012; Scott,Martens, and Tarr, 2013).D´emoulin et al. (2013) and Baker et al. (2017) tracked several ARs as theycrossed the solar disk. By applying a simple model of stationary upflow to theobserved ARs, they found that the long-term evolution of the persistent upflowsis consistent with the scenario of steady upflows projected onto the line of sight(LOS). They found no obvious dependence of the de-projected upflow velocityon the AR age. They were also able to determine the inclination angles of themagnetic-field lines that guide the upflows. Baker et al. (2017) found that theinclination angles are typically in the range of [0 ◦ ,40 ◦ ] and [-30 ◦ ,30 ◦ ] relative toto the local vertical direction for the trailing and leading polarities, respectively.These angles are roughly consistent with those inferred from magnetic-fieldextrapolations. The upflow velocities show no obvious difference between thetrailing and leading polarities and do not depend on the underlying photosphericmagnetic-field strength. They also found that transient coronal events such asCMEs, jets, and flares could lead to a variation of the upflow velocities ondifferent time scales. However, these transient changes normally do not affectthe quasi-steady behavior of the upflows. SOLA: upflow.tex; 9 February 2021; 2:31; p. 9 . Tian et al.
Figure 5.
Intensity image (shown with a reversed color table) and Dopplergram of theFe xii
Hinode /EIS (Harra et al., 2008).Reproduced by permission of the AAS.
These upflows first appear during the phase of flux emergence and generallypersist during the whole lifetimes of ARs. Harra et al. (2010) performed a detailedinvestigation of an emerging AR, and they witnessed the formation of theseupflows. As the AR expands, they observed a ring of blue shifts in coronalspectral lines around the edge of the emerging AR. The upflows clearly intensifyas more magnetic flux emerges. In a subsequent study, Harra et al. (2017) trackedthe evolution of a decaying AR during three solar rotations, and they found alarger area occupied by upflows and slightly higher upflow speeds as the ARevolves. Zangrilli and Poletto (2016) extended investigations of AR upflows onthe solar disk to the off-limb corona utilizing observations from the
UltravioletCoronagraph Spectrometer (UVCS: Kohl et al., 1995) onboard SOHO. Theirfindings demonstated that upflows persist over the lifetime of an AR spanningover four solar rotations.Through an examination of the coronal line profiles at AR boundaries, Haraet al. (2008) found that the profiles are generally asymmetric, with a weak yetnoticeable enhancement at the blue wings. These blueward asymmetries suggestthe presence of unresolved high-speed upflows. Later, detailed analyses of theseline profiles suggest that the coronal emission consists of at least two compo-nents: a primary component accounting for the nearly stationary backgroundemission and a secondary component associated with high-speed upflows (e.g.Tian et al., 2011b; Brooks and Warren, 2012). Figure 6 shows examples of suchline profiles, which also indicate a temperature dependence of the high-speedcomponent. Several methods have been introduced to separate the two compo-nents and characterize the weak high-speed component, including the techniqueof double-Gaussian fit (Peter, 2010; Bryans, Young, and Doschek, 2010; Brooks
SOLA: upflow.tex; 9 February 2021; 2:31; p. 10 pflows
Figure 6.
Double-Gaussian fits to several line profiles observed at an AR boundary (Brooksand Warren, 2012). The histograms represent the observed line profiles. In each panel, theprimary and secondary components are shown as the red and blue curves, respectively. Thegreen curve represents the total of the two components. Reproduced by permission of the AAS. and Warren, 2012; Doschek, 2012; Kitagawa and Yokoyama, 2015), red–blue(RB) asymmetry analysis (De Pontieu et al., 2009; De Pontieu and McIntosh,2010; Mart´ınez-Sykora et al., 2011; Tian et al., 2011b; Tian, McIntosh, and DePontieu, 2011) and RB-guided double-Gaussian fit (De Pontieu and McIntosh,2010; Tian et al., 2011b). Application of these methods to the observed coronalspectral-line profiles at AR boundaries mostly yielded speeds on the order of ≈
100 km s − , often in the range of 50 – 150 km s − , for these high-speed upflows.Occasionally the velocity may reach ≈
200 km s − . The primary component isoften also blueshifted, but by only ≈
10 km s − . The intensity ratio of the twocomponents is mostly less than 15 %, although at some locations it could reachmore than 30 %. The widths of the two components are often comparable, whichhas also been confirmed through a comparison with analysis of artificial profiles(Tian et al., 2011b). Since an observed spectral-line profile is the superpositionof the two components, an enhanced line broadening and a moderate blue shift( ≈
10 – 50 km s − ) are expected if a SGF is applied. This superposition also ex-plains the observed positive correlation of the blueward asymmetry with theblue shift and line broadening determined from a SGF (Tian et al., 2011b). It isworth mentioning that a blue-wing enhancement could in principle result fromthe superposition of multiple (more than two) emission components, each slightlyDoppler-shifted with respect to each other (Doschek, 2012). This scenario mightexplain some of the AR upflows but is not very likely for all of them, because atsome locations the blue-wing enhancement is so strong and far from the line corethat a multiple-component scenario is really not consistent with the observed lineprofiles. In addition, speeds of the counterparts of these upflows in coronal image SOLA: upflow.tex; 9 February 2021; 2:31; p. 11 . Tian et al. -600 -500 -400 -300-300-200-1000 Y ( a r cs e c ) Line Intensity (arbitrary unit) -600 -500 -400 -300-300-200-1000
Velocity (km/s) -30 -20 -10 0 10 20 30 -600 -500 -400 -300-300-200-1000
Width (km/s)
36 42 48 54 60 66 72 -600 -500 -400 -300-300-200-1000
RB Asymmetry [70-130 km/s] -0.21 -0.14 -0.07 0.00 0.07 0.14 0.21 -100 0 100 200-300-200-1000 Y ( a r cs e c ) -100 0 100 200-300-200-1000 -100 0 100 200-300-200-1000 -100 0 100 200-300-200-1000500 600 700 800X (arcsec)-300-200-1000 Y ( a r cs e c )
500 600 700 800X (arcsec)-300-200-1000 500 600 700 800X (arcsec)-300-200-1000 500 600 700 800X (arcsec)-300-200-1000
Figure 7.
Maps of the SGF parameters and profile asymmetry in the velocity interval of70 – 130 km s − (Tian et al., 2012b). The first, second, and third rows correspond to obser-vations of the same AR when it is located near the east limb, disk center, and west limb,respectively. Reproduced by permission of the AAS. sequences (see Section 3.2) can often be unambiguouly measured and they do notshow a continuous distribution.Although the blueward asymmetries are weak, they are definitely not causedby random noise or blends by other spectral lines. Tian et al. (2012b) havedemonstrated this by tracking an AR as it rotates from the east limb to thedisk center and then to the west limb. From Figure 7 we see a clear center-to-limb variation of the line parameters. Here the RB asymmetry is defined asthe difference of the two wing intensities integrated over the velocity intervalof 70 – 130 km s − divided by the peak intensity. When the AR is located nearthe disk center on 12 December 2007, we see prominent blue shifts, large linebroadenings, and blueward asymmetries at both boundaries. As the AR rotatesto the west limb on 15 December 2007, the profile asymmetries disappear atthe western boundary. This phenomenon can be understood if we consider thescenario of a high-speed field-aligned upflow superimposed on a nearly staticcoronal background. When the AR is close to the west limb, the magnetic-fieldlines at the western boundary are roughly perpendicular to the LOS so that theprojected speed of the field-aligned flow on the LOS direction is very small. Thisleads to a greatly reduced blueward asymmetry as well as a reduced Doppler shiftand line broadening. Similarly, when the AR is located near the east limb on 10December 2007, the blue shift, line broadening, and blueward asymmetry are allreduced at the eastern boundary because the field lines are roughly perpendicularto the LOS there. SOLA: upflow.tex; 9 February 2021; 2:31; p. 12 pflows
Figure 8.
Intensity images (shown with a reversed color table) and Dopplergrams of severalemission lines at an AR boundary (Warren et al., 2011). Contours of the Si vii (cid:48)(cid:48) × (cid:48)(cid:48) . Reproduced bypermission of the AAS. There is a clear dependence of the SGF velocity on the line-formation tem-perature at AR boundaries. Del Zanna (2008) found that the SGF blue shiftincreases from only a few km s − at 10 . K to ≈
30 km s − at 10 . K at an ARboundary. Tripathi et al. (2009) analyzed the EIS data at another AR boundary.They found an average red shift of a few km s − for the Si vii T ≈ ≈ − at10 . K to ≈
10 km s − at 10 . K. Actually, at AR boundaries both downflows andupflows could exist in the TR line Si vii T = 5.8 to 6.3 (Figure 8). The transition from red shiftsto blue shifts occurs at a temperature of log T ≈ − from the Fe xii SOLA: upflow.tex; 9 February 2021; 2:31; p. 13 . Tian et al.
Corona
Buffeting Convective Flows
Type-IISpiculesType-ISpicules
1) Heated Material (T>20kK) 2) RadiativelyCooling Material (T<1MK) λ I( λ ) λ Emission Line < 1MKBlue-Shifted ComponentBlue-Shifted Component RadiativelyCooling Component λ I( λ ) λ Emission Line > 1MK
Figure 9.
A cartoon showing the mass cycling between the chromosphere and corona(McIntosh et al., 2012). Reproduced by permission of the AAS. quence, consistent with the quasi-periodic blue-wing enhancements in the Fe xii lineprofiles. From the Si vii − , which are comparable tothe magnitude of red shifts at the locations of fans. From imaging observationsof AIA, Kamio et al. (2011) and McIntosh et al. (2012) have also identifiedquasi-periodic fast ( ≈
100 km s − ) outward-propagating disturbances in the hot193 ˚A channel and sporadic slow ( ≈
15 km s − ) downflows in cool channels suchas 131 ˚A. McIntosh et al. (2012) also found a mixture of upflows and downflowsfrom the 171 ˚A images, indicating that the dominant flow changes its directionaround a temperature of ≈ vii vii line profiles appears SOLA: upflow.tex; 9 February 2021; 2:31; p. 14 pflows to be smaller than that inferred from the simultaneously observed coronal lineprofiles, which is likely due to the fact that TR line profiles are complicatedby downflows. Unlike the two-component coronal line profiles, a TR line pro-file likely consists of three components: a fast upflow component generatedthrough chromospheric heating, a slow cooling downflow component, and a TRbackground component. The downflows have a much lower speed and possiblystronger emission compared to the upflows, which would result in net red shiftsof the TR lines when a SGF is performed. So similar to the quiet Sun and coronalholes, the temperature dependence of the SGF Doppler shift at AR boundariesis probably caused by the relative contributions of the different components atdifferent temperatures. A similar scenario of mass cycling or circulation has alsobeen proposed by Marsch et al. (2008) and Young, O’Dwyer, and Mason (2012).Sometimes the red-shifted downflows on the fan loops appear to be spatiallyuncorrelated with the blue-shifted upflows (as in one of the examples in Warrenet al., 2011). In these cases the mass-circulation picture might still be relevantif the upflows are connected to a distant AR (as in Boutry et al., 2012), i.e. thered shifts on the fans of the distant AR might be signatures of the circulation.In such cases the morphology of the fan loops and upflows is similar because ofthe general topology of the magnetic field at the active-region boundary (Bakeret al., 2009).
From EUV and X-ray imaging observations, we often find quasi-periodic upwardpropagating disturbances (PDs) along the fan-like structures at AR boundaries(e.g. De Moortel, Ireland, and Walsh, 2000; Sakao et al., 2007; Yuan and Nakari-akov, 2012). These PDs have a propagation speed of ≈
50 – 200 km s − , and theyoften recur at a time scale of 3 – 15 minutes. These PDs were widely interpretedas slow-mode magnetoacoustic waves propagating along the fans, mainly becausethe speed is comparable to the coronal sound speed and there is a correlationbetween the intensity and velocity perturbations (de Moortel, 2009; Wang et al.,2009; Wang, Ofman, and Davila, 2009; Nishizuka and Hara, 2011).The aforementioned blue shifts and these PDs are often found at the samelocations, indicating that they are closely related to each other. Based on thiscoincidence, the PDs have been suggested to be plasma upflows (Sakao et al.,2007; Harra et al., 2008). However, there appear to be some distinct differencesbetween them. First, around the year 2008, it was unclear why the blue shiftsare only ≈
10 – 50 km s − , which are significantly lower than the speeds of PDs. Aprojection effect has been proposed to explain the apparent discrepancy betweenthe LOS velocities of the blueshifts and the plane-of-sky (POS) velocities of thePDs. However, the blue shifts seldom exceed 50 km s − , regardless of AR locationon the solar disk. Second, the PDs are clearly quasi-periodic. However, it wasunclear whether the blue shifts showed a quasi-periodic variation around 2008.These discrepancies disappear if we compare the PDs with the high-speedsecondary components of coronal line profiles. From a coordinated observationwith XRT and EIS onboard Hinode , Tian, McIntosh, and De Pontieu (2011)found a clear correspondence between the fluctuations of the X-ray intensity
SOLA: upflow.tex; 9 February 2021; 2:31; p. 15 . Tian et al. -6-4-20246 F e X . Y ( a r cs e c ) Fe X 184.54 -6-4-20246 F e X II . Y ( a r cs e c ) Fe XII 195.12 -6-4-20246 F e X III . Y ( a r cs e c ) Fe XIII 202.04 -6-4-20246 F e X I V . Y ( a r cs e c ) Fe XIV 274.20 0 50 100 150Minutes after 18:37 on 2007/03/28 -6-4-20246 F e XV . Y ( a r cs e c ) Fe XV 284.16 -6-4-20246 S i V II . Y ( a r cs e c ) Si VII 275.35
Doppler shift (km/s) -30 -15 0 15 30
Intensity (% X 0.5)Inverted Doppler shift (km/s)Line width (km/s)Inverted RB asymmetry [60-120 km/s] (%)
Figure 10.
Quasi-periodic variation of the AR upflows (Tian et al., 2012a). Left: Temporalevolution of the Doppler shift of several emission lines. Right: Curves with different colorsrepresent the temporal changes of the line intensity, Doppler shift, line width, and profileasymmetry (RB asymmetry) averaged within the region between the two dashed lines in eachof the left panels. For a clearer illustration, the intensity and asymmetry curves are offset by3 and -4, respectively. Reproduced by permission of the AAS. and the blueward asymmetries of coronal spectral lines, suggesting that thePDs observed in XRT images are closely related to the secondary components.Through joint observations of SDO/AIA and
Hinode /EIS, Tian et al. (2011b)demonstrated that the velocity distributions and relative intensities are bothremarkably similar for the simultaneously observed PDs and secondary com-ponents. In addition, with sit-and-stare observations of EIS, De Pontieu andMcIntosh (2010) and Tian, McIntosh, and De Pontieu (2011) noticed that theblue shifts and blueward asymmetries often show quasi-periodic variations. Afollowing statistical study by Tian et al. (2012a) demonstrated that this behavioris quite common, and that the recurring time scale is similar to that of PDs. Soit appears that the PDs are related to the high-speed secondary components.Detailed analyses suggest that the blueward asymmetry as well as the inten-sity, blue shift and line width derived from a SGF all reveal correlated quasi-periodic variations. Figure 10 shows an example. The quasi-periodic enhance-ment of blue shift is clearly seen from the space–time diagrams. For the spectrallines formed at typical coronal temperatures of log T = 6.0 – 6.3 (Fe x xii xiii xiv xv SOLA: upflow.tex; 9 February 2021; 2:31; p. 16 pflows et al., 2012a). On occasions when there is only the nearly stationary background-emission component and no upflow, the Doppler shift is zero and the spectralprofile is symmetric. While when a high-speed upflow component coexists withthe background component, a SGF to the total emission profile will lead to anincrease in the line intensity and line width. In the meantime, the line profilereveals an enhancement in the blue wing and a SGF gives a small blue shift. Thisscenario naturally explains the in-phase variations of all line parameters, whichappears to favor the interpretation of intermittent plasma jets for the PDs. Ifthis new interpretation is correct, previous diagnostics of the coronal plasmabased on the wave interpretation of PDs would be significantly impacted.However, there is also observational evidence supporting the slow-wave inter-pretation. For instance, Verwichte et al. (2010) considered a scenario of propa-gating slow waves with a quasi-static background plasma component in the LOS.They were able to reproduce a blueward asymmetry in the line profiles, and a cor-relation of the blueward asymmetry with the line intensity and blue shift. Theyconcluded that the slow-wave interpretation is still valid for the PDs. However,their model predicted a double frequency (or half period) in the line width thatis not seen in EIS observations. There are also a couple of studies reporting atemperature-dependent PD velocity (Krishna Prasad, Banerjee, and Singh, 2012;Uritsky et al., 2013), which is a characteristic of slow-mode magneto-acousticwaves. However, several other studies did not reveal any obvious temperaturedependence (Tian et al., 2011b; Sharma et al., 2020). A statistical study byKiddie et al. (2012) showed that the temperature dependency appears to beclear only for PDs rooted in sunspots.After several years of debate, there is now a consensus that both wavesand flows exist at AR boundaries. For example, recent IRIS and AIA obser-vations have revealed signs of magneto-acoustic shock waves and jet-like flows,both of which appear to show a correspondence to the PDs (Bryans et al.,2016). As shown in Figure 11, the quasi-periodic sawtooth-like patterns in theMg ii h 1403 ˚A spectral profiles (evidence of magnetoacoustic shock waves, e.g.Rouppe van der Voort et al., 2003; Tian et al., 2014a; Skogsrud, Rouppe van derVoort, and De Pontieu, 2016), the Si iv blue shifts (signs of jet-like flows) and thePDs in 171 ˚A appear to be correlated. Thus both waves and flows may contributeto the signals of PDs. From EIS observations, the enhanced line width and blue-ward asymmetry are most obvious at the lower parts of fan structures, and theybecome absent at higher parts. This observational fact might be explained by aLOS projection effect or lower signals at larger heights. However, it could also beexplained as due to the dominance of slow-mode waves away from the footpointregions. Indeed, Nishizuka and Hara (2011) found an in-phase variation of thecoronal-line intensity and Doppler shift at larger heights, which was interpretedas an evidence of slow-mode waves. They also claimed that at lower heightsfast upflows obviously exist and may be driven by heating events around thefootpoints of loops.Attempts have been made to understand the coexistence of waves and flows.For instance, Ofman, Wang, and Davila (2012) performed 3D MHD modelingof hot ( ≈ SOLA: upflow.tex; 9 February 2021; 2:31; p. 17 . Tian et al.
Figure 11.
Contribution of PDs from both waves and flows (Bryans et al., 2016). Top panel: aspace–time diagram of the AIA 171 ˚A running difference for a virtual slit along the leg of a fanloop. The red horizontal line indicates the IRIS slit location. Second and third panels: temporalevolution of the Si iv ii h 1403 ˚A spectral line profiles. The black horizontal linesindicate the rest wavelengths of the two lines. Bottom panel: temporal evolution of the Mg ii h3Doppler shift (black) and the negative of the 171 ˚A intensity (normalized, red). Reproduced bypermission of the AAS. upflows can excite undamped slow-mode waves that propagate along the mag-netic loops. Their simulations also generated slow-mode shock-like wave trainswhen the driving pulses have a large amplitude. In a subsequent numericalinvestigation, Wang, Ofman, and Davila (2013) expanded this model to warmloops ( ≈ Baker et al. (2009) investigated the magnetic-field topology of an AR using alinear force-free field extrapolation method. From the reconstructed 3D coronalmagnetic field, they identified quasi-separatrix layers (QSLs), which are locations
SOLA: upflow.tex; 9 February 2021; 2:31; p. 18 pflows
Figure 12.
Hinode /EIS Fe xii
195 ˚A intensity (left) and Doppler velocity (middle) maps andphotospheric traces of QSLs in thick red lines overplotted on a grayscale Dopplergram (right).MDI magnetic field isocontours of ±
50 G are shown in white/black (middle). Field lines withcircles leave the computational box and are considered to be open field. Strong upflows (bluepatches in the middle panel, or dark patches in the right panel) over the positive polarity arealong open-field lines on the eastern edge of AR 10942 (Baker et al., 2009). The EIS imagescorrespond to a FOV of about 250 (cid:48)(cid:48) × (cid:48)(cid:48) . Reproduced by permission of the AAS. of strong gradients of magnetic connectivity. The strongest upflows appear tobe located in the vicinity of QSL sections over areas of strong field. Based onthis finding, the authors suggested that the upflows are driven by magneticreconnection at QSLs between closed field lines at AR cores and open-field linesor large-scale externally connected loops from AR boundaries. Edwards et al.(2016) analyzed EIS data of seven ARs and found that the upflows generally donot correspond to high-reaching loops or open-field structures predicted by theglobal potential-field source-surface model, but they also found that the upflowsoften coincide with the footprints of separatrix surfaces that are associated withcoronal null points. A good agreement between upflow regions and QSLs, eithertemporally or spatially or both, has also been found by Scott, Martens, and Tarr(2013), D´emoulin et al. (2013), Mandrini et al. (2015), and Baker et al. (2017).Similarly, D´emoulin et al. (2013) proposed that the upflows are driven by theupward pressure gradient after magnetic reconnection between the high-pressureAR loops and neighboring low-pressure loops. This scenario naturally explainsthe observational fact that stronger upflows are located closer to the AR cores.The steadiness of the upflows could be understood if we consider the scenario ofsuccessive reconnections, which leads to a superposition and thus averaging offlows with different velocities.Del Zanna et al. (2011) studied two ARs and found a clear association ofthe upflows with metric radio noise storms and large-scale open separatrix fieldlines. Based on this connection, they proposed that the upflows are driven byinterchange reconnection between magnetic loops in AR cores and adjacent open-field structures. In this scenario, continuous AR expansion leads to successivereconnection at coronal null points, which results in a strong pressure gradientthat drives the temperature-dependent plasma upflows. Bradshaw, Aulanier, andDel Zanna (2011) numerically simulated this scenario and found the development SOLA: upflow.tex; 9 February 2021; 2:31; p. 19 . Tian et al. of a rarefaction wave in the post-reconnection region. Their forward calculationyielded a ≈
10 – 50 km s − blueshift and clear temperature dependence of thevelocity magnitude, consistent with observational results based on a SGF (e.g.Del Zanna, 2008). On the other hand, a 3D data driven simulation of a similarmechanism by Galsgaard et al. (2015) failed to reproduce systematic upflows seenin EIS observations (Vanninathan et al., 2015). Instead, this model generatedmagneto-acoustic waves.Besides magnetic reconnection, AR expansion has also been considered as amechanism to drive the upflows. Murray et al. (2010) performed a 3D MHDsimulation, and they found that an upward acceleration of the coronal plasmais achieved when an AR expands horizontally in a unipolar background-fieldenvironment. They have managed to reproduce upflow speeds up to 45 km s − .Since AR expansion naturally leads to an intensification of the electric currentat the interface between closed- and open-field structures, which favors the oc-currence of magnetic reconnection, it is likely that reconnection and expansionboth contribute to the generation of the AR upflows. This was demonstratedwhen upflow velocities are significantly enhanced on the side of an AR adjacentto the open field of a nearby coronal hole. The slow rise and expansion of a fluxrope contained within the AR lead to stronger compression of the open fieldand the intensification of upflows hours before the flux rope erupts as a CME(Baker, van Driel-Gesztelyi, and Green, 2012). Using a similar approach, Harraet al. (2012) found that the upflows at the west and east sides of an emerging ARare dominated by reconnection jets and pressure-driven upflows, respectively.Hara et al. (2008) conjectured that high-speed upflows inferred from theblueward asymmetries of coronal line profiles result from impulsive heating atthe base of the corona, around the footpoints of AR loops. Based on an observedassociation of the episodic high-speed upflows with chromospheric activity, McIn-tosh and De Pontieu (2009b) suggested that these upflows represent discretemass-injection events produced by episodic local heating of the chromosphere.A similar scenario was also proposed based on the similar velocity distribu-tions of these high-speed coronal upflows and fast chromospheric spicules (DePontieu et al., 2009). In limb observations, some chromospheric spicules appearto be heated to TR and coronal temperatures, exhibiting as upward propagat-ing upflows in the AIA EUV passbands (De Pontieu et al., 2011). Based onthese observations, these authors proposed that the high-speed coronal upflowsare actually heating signatures during the upward propagation of spicule-likeplasma jets that are possibly driven by magnetic reconnection (e.g. De Pontieuet al., 2009; Samanta et al., 2019; Chen et al., 2019b) or other processes such asamplified magnetic tension (Mart´ınez-Sykora et al., 2017) in the chromosphere.They further pointed out that these episodic high-speed upflows generated atchromospheric heights play a key role in the mass and energy supply to thecorona.Although this scenario has received a lot of attention in the community, it hasalso been questioned by some studies. For instance, Klimchuk (2012) examined ascenario of spicule-supplied coronal plasma, which predicted a much larger blue-wing enhancement than observed. They concluded that spicules cannot provide SOLA: upflow.tex; 9 February 2021; 2:31; p. 20 pflows sufficient pre-heated plasma to fill the corona. Even if they did, additional heat-ing would be needed to maintain the high temperature as the plasma expandsupward and cools adiabatically. Klimchuk and Bradshaw (2014) and Bradshawand Klimchuk (2015) performed one-dimensional hydrodynamic simulations forthe idea of coronal heating by spicules, and they found that the synthesizedcoronal spectral-line profiles are distinctly different from the observed ones. Theyconcluded that impulsive heating events in the chromosphere cannot explain thebulk of coronal heating, although they may be responsible for heating of thechromospheric spicules to TR temperatures. With EIS observations of multiplespectral lines formed at different temperatures, Tripathi and Klimchuk (2013)derived the emission-measure distribution for the secondary component, andthey concluded that the emission measure is too small to support the proposalof coronal mass supply by chromospheric spicules. Using the density sensitiveline pair Fe xiv α line profile, the formation of which is highly complex. In contrast, a recentstudy by Polito et al. (2020) has found clear evidence of a correlation betweenthe coronal upflows and spectroscopic signatures in TR and chromospheric linesobserved with IRIS. The C ii iv ii k2 shows a positive asymmetry, which may be interpreted as signaturesof upflows.By examining the evolution of underlying photospheric magnetic field, Suet al. (2012) identified clear signatures of flux cancellation around one footpointof an AR loop. Fast jet-like upflows were found to initiate from the footpoint,and they are accompanied by transient brightenings in EUV images. However,no evident signature was found in images of the low chromosphere. They sug-gested that these upflows result from magnetic reconnection at the height of theupper chromosphere. Liu and Su (2014) found similar results, and they proposedthat the intermittent upflows are produced by reconnection between small-scaleemerging bipoles and pre-existing open-field structures at AR boundaries.Recent observations by the Parker Solar Probe (PSP: Fox et al., 2016) have re-vealed quasi-periodic Type-III radio bursts that are well correlated with coronalintensity variations at the footpoints of AR loops, suggesting electron accelera-tion during impulsive reconnection process around the loop footpoints (Cattellet al., 2020). It is unclear whether these radio bursts are signatures of thereconnection processes that generate the AR upflows, but Harra et al. (2021)have found that a radio noise storm during PSP Encounter 2 (between 31 Marchand 4 April 2019) likely originates in the expansion of the upflow region at theboundary of the only AR on the visible disk during that period.
SOLA: upflow.tex; 9 February 2021; 2:31; p. 21 . Tian et al.
Figure 13.
Source regions of the slow solar wind (Brooks, Ugarte-Urra, and Warren, 2015).From left to right: full-Sun images of the Fe xiii xiii
The upflows are generally found at the edges of ARs, where EUV spectrallines formed at typical coronal temperatures often show reduced emission. Theelectron density and temperature there are both lower than those in the coresof ARs (Doschek et al., 2008). Interestingly, coronal holes also show reducedcoronal emission, lower electron density, and lower temperature compared toneighboring quiet-Sun regions. These similarities between AR edges and coronalholes suggest that AR boundaries might also be source regions of the solar wind.Indeed, Kojima et al. (1999) found that low-speed solar-wind streams appear tobe associated with regions of large magnetic-flux expansion in the vicinity ofARs. This association was confirmed by Ko et al. (2006), who compared thecoronal electron temperature and abundances in the vicinity of an AR within-situ measurements of the slow solar wind.Magnetic-field extrapolations have been performed to understand the magnetic-field structures associated with the persistent upflows. Although a few inves-tigations showed that a portion of the upflows may be guided by large-scalemagnetic loops and eventually flow downward to the other footpoints of theloops (e.g. Boutry et al., 2012), most studies clearly demonstrated a closeassociation of at least some of these upflows with open-field lines originatingfrom AR boundaries (e.g. Marsch, Wiegelmann, and Xia, 2004; Sakao et al.,2007; Harra et al., 2008; Marsch et al., 2008; Baker et al., 2009). van Driel-Gesztelyi et al. (2012) performed both local magnetic-field extrapolations forindividual ARs and global potential-field source-surface modeling. They foundthat a part of the AR upflows are confined in closed-field structures, and thatthe other part of the upflows are associated with field lines extending to thesource surface from a coronal null point. This association strongly suggests thatthese AR upflows likely flow outward into interplanetary space and become partof the slow solar wind. However, Culhane et al. (2014) found that their studied
SOLA: upflow.tex; 9 February 2021; 2:31; p. 22 pflows
AR was completely covered by the closed field lines below the source surface,while the in-situ measurements, combined with the back-mapping technique,strongly suggest the origin of a slow-wind stream from the vicinity of this AR.To solve this apparent inconsistency, Mandrini et al. (2014) proposed a two-stepreconnection process. After the first reconnection between the expanding ARloops and neighboring large-scale loops, new large-scale loops connecting theAR upflow and a distant location are formed. These loops then reconnect withopen-field lines from a coronal hole around a coronal null point, releasing plasmaof the AR upflow into interplanetary space. Harra et al. (2017) have also shownthat the upflowing plasma could be released through interchange reconnectionif the upflow-hosting ARs are close to an open-field region.Chemical composition could be used to establish a link between some struc-tures in the interplanetary solar wind and their source regions in the solarcorona (e.g. Feldman and Widing, 2003; Song and Yao, 2020). Chemical elementswith a first ionization potential (FIP) below and above ≈
10 eV are often calledlow-FIP and high-FIP elements, respectively. It is well-known that the relativeabundance of a low-FIP element to a high-FIP element is generally enhancedin the slow solar wind relative to its photospheric value by a factor of threeto four. This factor is called the FIP bias. Spectroscopic observations at EUVwavelengths can be used to measure the FIP bias in the solar corona. Basedon
Hinode /EIS observations, Brooks and Warren (2011) found that the ratioof low-FIP Si and high-FIP S is enhanced by a factor of three to four at thelocations of AR upflows. A similar ratio was found in the in-situ measurementsof solar wind a few days later, providing evidence for a connection between theobserved solar-wind stream and the AR upflows. In a following study, Brooks,Ugarte-Urra, and Warren (2015) applied a similar technique to the EIS dataobtained through a full-Sun mosaic observation campaign and obtained a full-Sun composition (FIP bias) map. After comparing this composition map withthe simultaneously obtained full-Sun Dopplergram and the global magnetic-fieldtopology, they identified multiple locations of AR upflows as source regions ofthe slow solar wind (Figure 13).Using the technique of double-Gaussian fit, Brooks and Warren (2012) decom-posed the two components of the coronal line profiles observed at the boundariesof an AR. They found that the FIP bias for the high-speed upflow component isalso in the range of three to five, suggesting that the high-speed component mayalso contribute to the slow solar wind. The complexity of the outflow componentscontributing to the solar wind was also highlighted by Brooks et al. (2020), whoused
Hi-C (Rachmeler et al., 2019) high spatial resolution images at 172 ˚A totry to isolate the upflows cleanly from the cooler fan loops. Their results suggestthat the variability in solar-wind composition measurements might be explainedby activity in the source region.By applying the technique of Doppler dimming to the spectra obtained withSOHO/UVCS in the height range of 1.5 to 2.3 solar radii from the solar center,Zangrilli and Poletto (2012) and Zangrilli and Poletto (2016) identified coronaloutflows at the edges of ARs, and they found an increase of the outflow speedfrom ≈
50 km s − at 1.5 solar radii to ≈
150 km s − at 2.5 solar radii. Since thespeeds at 1.5 solar radii are generally smaller than those of the secondary com-ponents in EIS observations, it is possible that the fast upflows corresponding SOLA: upflow.tex; 9 February 2021; 2:31; p. 23 . Tian et al.
Figure 14.
Left: AIA 193 ˚A images taken during two solar rotations (Macneil et al., 2019).The solar wind source points are overplotted on the images. Right: time series of severalsolar-wind parameters, with associated source regions labeled and separated by vertical lines.In panels b or d the first two rows correspond to the solar-wind velocity and C /C ,respectively. The third row shows the sourcepoint longitude for each mapped data point. Theorange bar indicates longitudes corresponding to the vicinity of the CH. The fourth row showsthe magnetic-field polarity of the radial component of the interplanetary magnetic field (circles)and the corresponding PFSS magnetic-field polarity for each mapped data point (black line).Reproduced by permission of the AAS. to the secondary components mostly supply mass to the corona and do notdirectly escape to the interplanetary space. Instead, the primary components inEIS observations, with a blue shift of about 10 km s − (e.g. Tian et al., 2011b),may be the major source of the outflows detected by UVCS. It is also possiblethat what UVCS measured is an average speed, meaning that the weak andtransient high-speed upflows may have been smeared out over the course oftemporal or spatial sampling.Longer-term measurements are also helpful for understanding the linkagebetween the upflows and the solar wind. By comparing the synoptic photo-spheric magnetograms and the time sequences of solar-wind parameters, Heet al. (2010) concluded that the upflowing plasma at an AR boundary mayevolve into an intermediate-speed solar-wind stream observed near the Earth. SOLA: upflow.tex; 9 February 2021; 2:31; p. 24 pflows
This solar-wind stream has a speed of about 400 km s − , and an intermediatetemperature and density if comparing to the typical fast and slow solar wind. Inanother study by Janardhan, Tripathi, and Mason (2008), the authors managedto establish a link between an extremely slow ( (cid:54)
300 km s − ) solar-wind streammeasured in interplanetary space and an evolving low-emission coronal regionlocated near an AR. They claimed that interchange reconnection between theAR loops and surrounding open fields leads to a reduction in the area of theopen-field region, which would result in a smaller solar-wind velocity. With thehelp of the potential-field source-surface (PFSS) model, Liewer, Neugebauer, andZurbuchen (2004) ballistically mapped several solar-wind streams back to thesource surface and then to the photosphere. They found that these solar-windstreams generally can be traced back to dark areas in EUV and soft X-ray imagesat the edges of ARs. Fazakerley, Harra, and van Driel-Gesztelyi (2016) trackeda Carrington rotation that had coronal holes and active regions. The solar windwas linked back to the features on the Sun, using ballistic back mapping andPFSS modelling. They found short periods of enhanced-velocity solar wind atthe boundary of slow and fast wind streams, which are related to ARs that arelocated beside coronal holes. The neighbours of ARs are important in termsof how the solar wind is created. Figure 14 shows an example of connectingremote sensing data that are showing regions of upflows in a coronal hole andan AR, with the solar-wind measurements (Macneil et al., 2019). These authorsexamined two solar rotations, one with the quiet Sun and coronal hole only,and during the next rotation an AR has emerged, making it ideal to comparethe two different situations. Their observations indicate that the features thatemerge in the AR-associated wind are consistent with an increased occurrenceof interchange reconnection during solar-wind production, compared with theinitial quiet-Sun case. Back-mapping the solar wind to the source surface totrace solar wind source regions has also been trialed on Hinode and ACE obser-vations as a science preparation for
Solar Orbiter (Stansby et al., 2020), and theauthors discussed some of the difficulties involved in interpreting the compositionmeasurements. A direct link between solar-wind streams and their coronal sourceregions may also be achieved in a different way by using the Wang – Sheeley –Arge (WSA) model (Wang and Sheeley, 1990; Arge and Pizzo, 2000). Using theWSA model, Slemzin et al. (2013) have also identified locations of AR upflowsas the source regions of slow solar-wind streams.Tracing the solar wind back to the Sun from 1999 to 2008, Fu et al. (2015)classified the solar wind by its source-region types with the EUV images andphotospheric magnetograms taken by SOHO. They found that about half of thesolar-wind streams come from AR regions during the solar maximum. The chargestates, FIP bias and helium abundance are significantly different for the solarwind coming from ARs, the quiet Sun, and coronal-hole regions (Fu et al., 2017,2018). Their results further indicate that the reconnection between the openmagnetic-field lines and closed loops may play a major role in the generation ofthe AR-associated wind.
SOLA: upflow.tex; 9 February 2021; 2:31; p. 25 . Tian et al.
4. Upflows from CME-Induced Dimmings
When coronal mass ejections (CMEs) leave the Sun, often we see reduced emis-sion of the corona, most noticeably at EUV and soft X-ray wavelengths (e.g.Hudson, Acton, and Freeland, 1996; Sterling and Hudson, 1997; Dissauer et al.,2018, 2019). This phenomenon is called coronal dimming, and a dimming regionis sometimes called a transient coronal hole. Statistical studies show that morethan half of the frontside CMEs are associated with dimmings (Reinard andBiesecker, 2008; Bewsher, Harrison, and Brown, 2008). There are at least twotypes of coronal dimmings: core (or twin) dimmings and secondary dimmings.Core dimmings usually refer to small regions of greatly reduced coronal emis-sion near the eruption sites of CMEs. They are generally believed to markthe footpoints of ejected flux ropes, and they could last for more than tenhours (e.g. Hudson, Acton, and Freeland, 1996; Vanninathan et al., 2018; Chenet al., 2019a; Xing, Cheng, and Ding, 2020). Secondary dimmings, which mayresult from stretching of the magnetic-field lines or reconnection of the eruptingmagnetic field with the surrounding structures, often show a less prominentintensity decrease and a quicker recovery (e.g. Thompson et al., 2000; Mandriniet al., 2007; Zheng, Chen, and Wang, 2016; Vanninathan et al., 2018). Note thatsome dimmings are likely not associated with CMEs, e.g. long-duration remotedimmings associated with confined circular-ribbon flares (Zhang and Zheng,2020) and dimmings at the peripheries of emerging flux regions (Zhang et al.,2012). Here we focus on CME-induced dimmings. Spectroscopic observationsoften reveal an obvious density decrease in these dimming regions (Harrison andLyons, 2000; Harrison et al., 2003; Tian et al., 2012b). In addition, the peaktemperature of the differential-emission-measure curve appears not to changewhen a dimming occurs (Tian et al., 2012a). These results demonstrate thatcoronal dimming is mainly due to mass loss rather than temperature change.Signatures of coronal dimming are also present in the Sun-as-a-star spectra takenby the
EUV Variability Experiment (EVE: Woods et al., 2012) onboard SDO(Mason et al., 2014).There are only a few spectroscopic investigations of coronal dimmings. UsingCDS observations, Harra and Sterling (2001) reported an obvious blue shiftin coronal and TR lines in coronal dimming regions. At least a few coronaldimming events have been caught by the EIS spectrograph. These observationsoften reveal a prominent blue shift of coronal emission lines in dimming regions,normally in the range of ≈
10 – 50 km s − , with an average of about 20 km s − when a SGF is performed (Harra et al., 2007, 2011). These blue shifts appear tobe very similar to those at AR boundaries, and they likely represent systematicplasma upflows along the field lines opened by CMEs. The strongest blue shiftsare found at the footpoints of disrupted coronal loops (Harra et al., 2007; Attrillet al., 2010). The blue shift appears to be positively correlated with the de-pression of coronal emission and magnetic-field strength in the photosphere (Jinet al., 2009). In addition, the spectral-line profiles are found to be broadenedcompared to the line profiles in the pre-eruption phase and outside dimmingregions. The large line width has been interpreted as a result of amplified Alfv´en SOLA: upflow.tex; 9 February 2021; 2:31; p. 26 pflows
Figure 15.
Maps of the peak intensity, Doppler shift, and line width derived from a SGF,and RB asymmetry in the velocity interval of 70 – 130 km s − for the Fe xiii (cid:48)(cid:48) × (cid:48)(cid:48) . The same dataset has alsobeen analyzed by Harra et al. (2007) and Jin et al. (2009). Reproduced by permission of theAAS. wave amplitude (McIntosh, 2009) or inhomogeneity of flow velocities along theLOS (Dolla and Zhukov, 2011).Similarly, the coronal line profiles have been found to be asymmetric indimming regions. Using the technique of RB-asymmetry analysis introducedby De Pontieu et al. (2009), McIntosh, De Pontieu, and Leamon (2010) founda couple of small patches in the dimming region behind an erupted halo CMEwhere the Fe xiii − . The intensity ratio of the two components is often about SOLA: upflow.tex; 9 February 2021; 2:31; p. 27 . Tian et al. −400 −200 0 200velocity [km/s]500.1000.1500.2000.2500. i n t e n s it y [ e r g / c m / s / s r / Å ] x = 557.2"y = 43.9" v = −1.6 ± = −134.5 ± −400 −200 0 200velocity [km/s]1000.2000.3000. x = 560.2"y = 43.9" v = 7.2 ± = −92.9 ± −400 −200 0 200velocity [km/s]1000.2000.3000.4000. x = 563.2"y = 43.9" v = 8.3 ± = −98.3 ± −400 −200 0 200velocity [km/s]1000.2000.3000.4000.5000. x = 566.2"y = 43.9" v = 4.5 ± = −78.3 ± −300 −200 −100 0 100 200velocity [km/s]5000.10000.15000. i n t e n s it y [ e r g / c m / s / s r / Å ] x = 557.2"y = 43.7" v = 1.4 ± = −127.4 ± −300 −200 −100 0 100 200velocity [km/s]5000.10000.15000.20000. x = 560.2"y = 43.7" v = −2.8 ± = −130.5 ± −300 −200 −100 0 100 200velocity [km/s]5000.10000.15000.20000. x = 563.2"y = 43.7" v = 1.0 ± = −129.7 ± −300 −200 −100 0 100 200velocity [km/s]5000.10000.15000.20000. x = 566.2"y = 43.7" v = 4.4 ± = −115.1 ± Figure 16.
Line profiles of Fe xiii xv ≈
20 km s − blue shift and enhanced linewidth derived from a SGF are caused at least partly by the superposition of thetwo components. So it appears clear that a fraction of the coronal plasma indimming regions flows upward at a speed of the order of ≈
100 km s − . RecentEUV-imaging observations with AIA have clearly revealed upflows with speedsof ≈
70 – 140 km s − from a CME footpoint in a dimming region (L¨orinˇc´ık et al.,2021), confirming the finding of Tian et al. (2012b).Similar blueward asymmetry and a high-speed upflow component have beencommonly observed in other dimming events. Chen, Ding, and Chen (2010)reported an expanding dimming event with a leading edge/ridge of enhancedline width. Tian et al. (2012b) examined the asymmetry of the coronal spectraand found clear blueward asymmetries along the whole ridge and in the obviousdimming regions behind it. In a recent study, Veronig et al. (2019) found thatat the growing dimming border the upflow component is so strong that thespectra even show distinct double components (Figure 16). At some pixels theupflow component even dominates over the nearly stationary component. Suchline profiles can be decomposed with a high degree of accuracy through a double-Gaussian fit.What is the significance of these upflows? Dimming regions are often calledtransient coronal holes, as the coronal magnetic-field lines are transiently openedby the CME ejecta. So a large fraction of these upflows likely flow outward alongthe opened field lines and eventually become solar-wind streams following theCME ejecta. In this sense, a dimming region should be regarded as a sourceregion of the solar wind, similar to AR boundaries and coronal holes (McIn-tosh, 2009; Tian et al., 2012b). The enhanced nonthermal broadening may becontributed by both the superposition of different emission components andthe growing Alfv´en wave amplitude in the open-field and low-density coronalenvironment. The momentum flux resulting from the high-speed upflows mayact as a secondary momentum source, and thus it may impact the kinematicsof the associated CMEs (McIntosh, De Pontieu, and Leamon, 2010). It is also SOLA: upflow.tex; 9 February 2021; 2:31; p. 28 pflows possible that these upflows continuously feed the associated CMEs (Harra et al.,2007). Using the density sensitive line pair Fe xiii × cm − for the upflow com-ponent. They then estimated a mass-loss rate due to the upflows, which is about8 × g s − and comparable to the mass-increase rate of the associated CMEinferred from coronagraph observations. Thus, the upflows may become part ofthe CME. In addition, since a dimming region will eventually recover to thepre-eruption state, part of these upflows might provide mass to refill the coronaafter the eruption of CMEs (Tian et al., 2012b).Besides these prevalent high-speed upflows that appear to show no obvioustemperature dependence, clear temperature-dependent upflows associated withdimmings have also been identified. For instance, Imada et al. (2007) identifiedan upflow with a speed increasing from ≈
15 km s − at 10 . K to ≈
160 km s − at 10 . K at the edge of a dimming region. With the RB-asymmetry technique,Tian et al. (2012b) detected several more temperature-dependent upflows, butall in very small areas immediately outside the (deepest) dimming regions.These temperature-dependent upflows are likely evaporation flows induced byinteraction between the stretching field lines in dimming regions and adjacentmagnetic-field structures.Dimmings have been linked to magnetic-cloud materials. Harra et al. (2011)used upflow measurements to determine accurately where the dimming occurredthrough a velocity-difference method. They used this to determine the maximumamount of magnetic field that could end up in a magnetic cloud. The magnetic-cloud measurements and modelling determined that the magnetic flux in thedimming region is consistent with the magnetic-cloud data.
5. Summary and Future Perspectives
In the past two to three decades, EUV/FUV spectroscopic observations havelargely changed our view of the upper solar atmosphere. With these dedicatedobservations, highly structured upflows and downflows with varying magnitudesat different temperatures have been identified in various regions on the Sun.On large scales these flows appear to be quasi-steady and long-lasting, possiblyindicating a continuous global plasma circulation (Marsch et al., 2008). System-atic upflows in the upper TR and corona have been commonly believed to besignatures of the nascent solar wind. High-cadence observations and detailedanalyses of the spectral-line profiles suggest that there are intermittent high-speed upflow components with a wide range of temperatures and a relatively slowcooling downflow component mostly at TR temperatures. These observationsindicate that the upper atmosphere is not a static and magnetically stratifiedlayer, but rather a dynamic interface revealing a continuous mass cycling betweenthe chromosphere and corona/solar wind.We have presented an extended overview of plasma flows observed by theEUV/FUV spectrographs of
Hinode /EIS, IRIS, SOHO/SUMER, and SOHO/CDS,with an emphasis on the upflows seen in the network structures of the quiet Sun
SOLA: upflow.tex; 9 February 2021; 2:31; p. 29 . Tian et al. and coronal holes, boundaries of ARs, and CME-induced dimming regions. De-spite significant advances in the research of these plasma flows, several importantunresolved issues remain and require further investigations.For the quiet Sun and coronal holes, there is still no consensus on the phys-ical mechanisms behind the observed temperature dependence of the Dopplershift. It is commonly believed that this temperature dependency is a conse-quence of coronal heating (Peter, Gudiksen, and Nordlund, 2006; Hansteenet al., 2010). So more realistic numerical simulations of coronal heating shouldbe performed to better understand this temperature dependency. In addition,the exact temperature at which the net Doppler shift changes sign is still un-clear, mainly because there are no strong emission lines formed between theredshift-dominated temperatures of log T (cid:54) . T (cid:62) . vii
465 ˚A line formed around log T = 5 . Spectral Imaging of the Coronal Environment (SPICE:SPICE Consortium et al., 2020) instrument onboard the newly launched
SolarOrbiter mission (M¨uller et al., 2020) has the capability of sampling several strongemission lines formed at both TR and coronal temperatures at a high spatialand temporal resolution. Combined observations between SPICE and the 4-m
Daniel K. Inouye Solar Telescope (DKIST) will likely reveal new insights intothe generation and propagation of the upflows/waves as well as their role in theenergization of the coronal plasma. Having a mission such as
Solar Orbiter in anorbit away from the Earth’s orbit will also allow 3D spectroscopy by combining
Solar Orbiter /SPICE with the Earth-orbiting
Hinode /EIS and IRIS missions. Amain goal of the
Solar Orbiter mission is to trace the interplanetary solar-windstreams back to their sources in the corona more accurately and routinely bycombining imaging and spectroscopic observations with in-situ measurements.FIP bias measurements both spectroscopically and in situ will form a significantpart of this work. New insight could also be obtained through observations with afuture spectrometer, the
Solar-C EUV High-Throughput Spectroscopic Telescope
SOLA: upflow.tex; 9 February 2021; 2:31; p. 30 pflows (EUVST, Shimizu et al., 2019), which has been approved by JAXA.
Solar-C /EUVST will provide near continuous spectral measurements throughout thesolar atmosphere at similar spatial resolutions. The spatial resolution will bearound seven times better than that of
Hinode /EIS, and the temperature cover-age allows us to observe seamlessly from the chromosphere ( ≈ ≈ ≈
10 million kelvin). Normallymagnetic-field extrapolations are required to guide our understanding of thecorona–solar-wind connection. With recently developed promising techniquesof coronal magnetic-field measurements (Yang et al., 2020a,b; Li et al., 2015,2016; Si et al., 2020; Landi et al., 2020), magnetic-field extrapolations will likelyimprove and thus help establish a more accurate connection between the coronalupflows and interplanetary solar-wind streams. In addition, the coronal magneticfield measurements from DKIST will allow us to fully understand the coronalmagnetic field rather than rely completely on models. These new facilities willrevolutionize our understanding of the energy transport in the solar atmosphere(Velli et al., 2020).Formation mechanisms of the upflows in CME-induced dimming regions, andtheir potential role in solar-wind formation and impact on CME evolution, arestill poorly understood. This is mainly due to the infrequent spectroscopic obser-vations of dimmings because of the small fields of view and lower cadence of theslit scans. Possibly, future instrumentation should aim at full-disk spectroscopicimaging with a wide temperature coverage to catch a large number of dimmingevents. In the
Solar Orbiter era, we also expect that the linkage between dimmingregions and solar-wind streams will be made more easily than before.
Acknowledgments
H. Tian is supported by NSFC grants 11825301 and 11790304. Thework of D. H. Brooks was performed under contract to the Naval Research Laboratory andwas funded by the NASA Hinode program. D. Baker is funded under STFC consolidated grantnumber ST/S000240/1. L. Xia is supported by NSFC grants 41974201 and 41627806. H. Tianacknowledges support from the UCL-PKU strategic partner funds during his visit to MSSL.This article is based upon the AAS/SPD Karen Harvey Prize Lecture of 2020, the presentationfile of which is available at http://spd.aas.org/prizes/harvey/previous.
Disclosure of Potential Conflicts of Interest
The authors declare that they have noconflicts of interest.
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