Energy release and transport in solar flares & CMEs

Transcription

1 UCL DEPARTMENT OF SPACE AND CLIMATE PHYSICS MULLARD SPACE SCIENCE LABORATORY Energy release and transport in solar flares & CMEs Sarah Matthews

2 Red RHESSI 6-12 kev, blue kev, gold images TRACE 195A

3 Flare - definition and classification A solar flare is a sudden release of energy during which via magnetic reconnection free magnetic energy is converted to kinetic energy of fast particles, mass motions, and radiation across the entire electromagnetic spectrum. Energy released up to J or erg in the largest solar flares. Many more much smaller flare-like events occur, with energies of as small as J nanoflares, micro-flares etc. GOES soft X-ray classification is most common these days. The flux in the 1 8 Å = nm range is recorded by this scheme: B 10-7 W/m 2 C 10-6 W/m 2 M 10-5 W/m 2 X 10-4 W/m 2 - so an M5 flare has flux of W/m 2. GOES = Geostationary Environmental Operational Satellites they are continuously recording solar X-ray emission.

4 Flare time profile Impulsive phase primary energy release hard X-rays (10s of kev) white light, UV, µwaves - broad spectrum duration < few minutes intermittent and bursty time profile, 100ms energy injection: - few tenths of the total flare energy released (up to ergs) - significant role for non-thermal electrons Gradual phase - response to input thermal emission (kt ~0.1-1 kev) rise time ~ minutes Neupert effect (Neupert, 1968): F SXR (t) t t 0 F HXR (t ' )dt ' Soft X-rays mainly originate from plasma heated by the accumulated energy deposited by accelerated electrons from flare start.

5 Flare frequency spectrum The number of flares falls off with increasing power as a flat power law with a slope of ~ -1.8 (SXR, EUV, microwave, HXR bursts, optical flares) (e.g. Drake, 1971; Dennis, 1985, Hudson, 1991) <1-hour dataset! Aschwanden et al, (2000) α=-1.8 dn/dw=a W -α (ergs s) -1 the normalisation factor A varies with the level of activity (Kreplin et al, 1977, Wagner, 1988) Recently α values were found between 1.5 and 2.6. Relevance for coronal heating: If α<2, smaller flares do not contribute enough to heat the corona.

6 Light curve of typical superflares. H Maehara et al. Nature 000, 1-4 (2012) doi: /nature11063

7 Comparison between the occurrence frequency of superflares on G-type stars and those of solar flares. Astronomical Society of Japan Kazunari Shibata et al. Publ Astron Soc Jpn 2013;65:49

8 Magnetic indicators of imminent flares? Intuition: Complex, rapidly evolving active regions have the highest probability of flaring More quantitative methods? Past X-ray activity (Bayesian stats) Wheatland (2004) moderately successful Magnetic field stats and variations No consistent picture emerged (Leka & Barnes, 2006) >70% data flare quiet at C1.0 level, so by doing nothing you get >70% success rate.

9 Prior flaring Falconer et al. (2012) Red prior flaring; Blue no prior flaring

10 High-gradient + strong-field - flaring Schrijver (2007) introduced R (new metric): summed unsigned B of high-gradient strong-field polarity inversion line. (overlap of +ve or -ve B > 150 Mx cm -2 Kernels: 6 x6 SOHO/MDI Φ= R x 2.2x10 16 Mx Schrijver, 2007 : M&X flares Forecast of major flare within 24 hours: R 2x10 21 Mx (logr 4.8), probability 1 R Mx (logr 2.8), probability 0 These features are characteristics of new flux emergence in highly non-potential state. Strong R, Φ & big flare correla2on

11 Origin and storage of free magnetic energy Origin: Magnetic flux emerges twisted, i.e. in a non-potential state, from the solar interior. Twist may keep propagating from below via torsional Alfvén waves. Surface flows and magnetic footpoint shuffling shear and entangle field lines. Storage: Magnetic free energy (above the energy of the potential state) is stored relatively low in an AR 20 Mm above the photosphere and may mainly be concentrated along the magnetic inversion line in the filament channel. Hinode/SOT magnetogram Integrated electric currents in an AR before (a) and after an X-class flare. Note their organization into an apparent flux-rope structure (Schrijver et al., 2008)

12 Twisted flux emergence and X-class flare Hinode/SOT G-band Upper photosphere WLF < 100 km height Isobe et al, 2007 Hinode/SOT Ca II H-line Chromosphere Kubo et al., 2007

13 How is the energy released? Magnetic reconnection is a topological restructuring of a magnetic field caused by change in the connectivity of its field lines. It allows the release of stored magnetic energy (dominant free energy in plasma) Evidence of reconnection has now (we believe) been seen Observation of the energy release site remains controversial Many pieces of indirect evidence in solar observations.

14 CSHKP model for eruptive solar flares (Carmichael 1964; Sturrock 1966; Hirayama 1974; Kopp & Pneuman 1976) prereconnec tion post reconnection prereconnec tion Filament Accelerated electons gyrate along magnetic field lines emitting gyrosynchrotron radiation Collisions in the dense chromosphere emits bremstrahlung observed in hard X-rays (> 20 kev). Electrons impulsively heat the chromosphere leading to optical and UV emission. Heated chromospheric plasma expands upward, increasing ρ and T in the reconnected coronal loops. post reconnection Adapted from Shibata (1998)

15 Confined flares - quadrupolar reconnection Reconnection happens at nullpoints (X-point) at separatrices and their intersection, the separator at quasi-separatrix layers (QSLs) Signatures: four flare kernels/ribbons at the footpoints of reconnected loops ribbons are in the vicinity of drastic field line connectivity changes. Priest & Forbes, 2000

16 Reconnection along QSL Along QSLs field line mapping is continuous but shows steep gradients. Reconnection along QSL does NOT break and reconnect field lines, but field lines may slip across each other, as shown in MHD simulations. The movie shows a case for sliprunning reconnection observed with Hinode/XRT. (Aulanier et al., 2007)

17 Su et al., 2013

18 Rising reconnection region, cooling loops Hinode/EIS: Cooler loops lie below the hotter loops since the lower ones were formed before the higher ones

19 Shrinking reconnected loops First examples were found in Yohkoh/SXT data McKenzie & Hudson (1999) TRACE EUV Shrinkage of newly reconnected cusped loops driven by the magnetic tension force. Patchy and intermittent reconnection process.

20 Patchy and intermittent reconnection Asai et al., 2004 Start of these downflows are associated with non-thermal HXR emission (RHESSI) and microwave bursts (NoRH).

21 Su et al., 2013

22 In 3-D Musset et al., 2015 Janvier et al Zharkov et al., 2011

23 How & where are particles accelerated? New insights from 3D models: Coronal X-ray emission overlies current ribbons in the photosphere New >50 kev HXR source appears in association with increased photospheric current > clear link between particle acceleration and reconnecting current sheets Musset et al., 2015

24 Hard X-rays Produced by electron-proton bremsstrahlung from electrons >15 kev 1 I( ε) = 4πR n V F (E)Q(ε, E)dE 2 ε Thermal bremsstrahlung: E electron ~ E target and spectrum F(ε) ~ e -ε/kt Non-thermal bremsstrahlung: E electron >> E target and spectrum F(ε) ~ ε -γ Very inefficient, ~ 10-5 electron energy radiated as X-rays In impulsive phase, HXR spectrum can be fitted by a hot (20 MK) or superhot (~60 MK) thermal component plus a power law.

25 Radio emission frequency Isliker & Benz 1994 frequency time time Upward and downward-going beams are observed, occurring at peak time of HXR emission. Metric and decimetric Type III bursts are often plasma radiation produced by electron beams (from Langmuir waves at f ~ 9 n e ).

26 Impact of accelerated particles Precipitation of energetic electrons and ions from the coronal acceleration site in the dense chromosphere heating overpressure plasma upflows (chromospheric evaporation)+ bright flare ribbons + HXR & γ-ray sources along the ribbons.

27 Gamma-ray footpoints HXRs and gamma-ray lines have similar time profiles, implying related acceleration, but ion signature is in different location from electron signature. Neutrons produced by energecc ions (10s of MeV/nucleon. Capture line predicted to form within 500 km of neutron produccon site. Observed offset from HXRs, ~10000 km Hurford et al., 2006

28 Flare quakes Seismic disturbances seen during the impulsive phase of a small number of flares. First detected by Kosovichev & Zharkova (1996) Approximately co-spatial with HXR and WL emission. How does the energy get so deep? 15 scale heights! Kosovichev & Zharkova, 1998

29 Kosovichev & Zharkova, 1998

30 Flare related magnetic field changes Magnetic reversals are seen in some flares, spatially and temporally correlated with HXR sources flare related changes to the line profile? Also sudden and permanent changes in the longitudinal field of ~ 10% ( G) (Sudol & Harvey (2005). Some are co-spatial with flare ribbons/ kernels and propagate at flare ribbon speed. Johnstone et al., 2009

31 Magnetic re-structuring and Alfvén waves Alfven waves heat chromosphere & drive evaporation (Reep & Russell, 2016)! Increased line widths, coronal upflows, lack of HXR signature Restructuring causes increase in horizontal B/ change in tilt Increased photospheric currents Russell et al., 2016 Reep & Russell, 2016

32 Evidence? Broadening of chromospheric lines/ absence of co-spatial HXR emission v D =134.7 km s -1 => wave energy flux ~ erg cm -2 s -1 Increased j at SQ location Mg II k Matthews et al., 2015 Sharykin et al., 2015

33 CMEs observable change in coronal structure that occurs on a timescale of a few minutes to several hours, and involves the appearance and outward motion of a new, discrete, bright, white light feature in the coronagraph field of view. (Hundhausen, 1986)

34 First observations Eclipse drawing (Tempel) 18 Jul 1860 Skylab 10 June 1973 (MacQueen et al. 1974)

35 Eddy 1974

36 Morphology 3 part structure: Bright fontal loop (overlying arcade) Dark cavity (flux rope) Bright core (filament/ prominence) ~30% show this structure

37 Jets, Halos and partial halos Narrow jet-like structures ~20 Partial and full halos ( ) Directed along the Sun-Earth line. 10% of all CMEs are halos; 4% full.

38 Thomson scattering CMEs are best observed in WL Thomson scattering of photospheric light by coronal electrons Depends on density of scattering electrons and angle between incident radiation direction and the l.o.s. Scattering is strongest in the plane of the sky, i.e. limb CMEs are favoured. Vourlidas & Howard, 2006

39 Different perspectives Single viewing perspectives can be misleading and give skewed perception of angular widths and other properties. 70 separation of STEREO A and B from LASCO. STEREO A LASCO STEREO B

40 Radio signatures CME speeds > v A drive a shock ahead of the CME which can accelerate electrons -> Langmuir waves -> Type II radio bursts Bastian et al., 2001

41 Properties Contain coronal material ~ few MK Cool prominence material in the core ~ 8000 K Compressed sheath behind the shock higher T and n e. Cavity lower density -> higher magnetic field for pressure balance. Virtually no measurements

42 Angular width Extent of PA in plane of the sky => only get accurate results for limb CMEs. Widths range from Average computed for < 120.

43 Speed and Acceleration CMEs start from rest. Driving force close to the Sun, interaction with solar wind slows the CME. Speeds measured from fits to H-t plots Fast CMEs decelerate, slow accelerate. a = (V-466) LASCO (Gopalswamy 2010)

44 Kinematics Mass: estimate number of electrons needed in plane of the sky to produce observed brightness. M = (B obs /B e (θ))x1.97x10-24 g; B e (θ) brightness of a single electron at angle θ from the plane of the sky (Vourlidas et al., 2000). Speed fits to height-time measurements in the plane of the sky (more reliable for limb than halo) V = W (W angular width)(gopalswamy et al. (2009)) Wider CMEs are generally faster and more massive.

45 Mass and KE Faster and wider CMEs have higher KE KE ranges from < to > Distributions become more symmetric if only limb CMEs are considered, and median values increase: 1.3x10 15 g and 1.6x10 30 ergs. Gopalswamy 2010

46 Statistical properties Table from Webb and Howard, 2012

47 Flares and CMEs Both arise due to reconfiguration and subsequent release of energy in the coronal magnetic field. Plenty of flares which are not accompanied by CMEs, or are failed eruptions. CMEs can also occur without flares, although many models predict flares as part of the eruption process. Correlation increases with flare size. 100% for >X3.0 (Yashiro et al., 2005)

48 On-disk signatures Flare ribbons and arcades Coronal dimming EUV waves Patsourakos et al., 2009

49 Occurrence rate Rate varies from <0.5/day and minimum to > 6 at max. Can be > 10/day Significant differences between catalogues automatic vs manual; observer bias; sensitivity. Robbrecht et al., 2009

50 Solar cycle variation CME occurrence rate follows the solar cycle in phase and amplitude Linear relationship between CME rate and sunspot number (Webb & Howard, 1994; Robbrecht et al., 2009) But correlation is higher during rise and declining phases than during max. Speeds are generally higher at max. Robbrecht et al. 2009

51 Latitudinal distribution Gopalswamy 2010

52 Halos

53 Halo speeds

54 Shocks Fast CMEs drive shocks Shock accelerates electrons, which produce Langmuir waves. Converted to Type II radio emission at plasma frequency. Drifts to decreasing frequency -> density. Direct imaging of shock sheath only really just been recognized cloud of electrons ahead of the CME. Rouillard et al., 2012

55 CMEs and SEPs Gradual SEP events strongly associated with CMEs Impulsive flares Also hybrid events Sources in West magnetically well connected to Earth But CME flanks are wide. Reames (1999)

56 CMEs and SEPs STEREO is showing that SEP events have up to 360 impact. Dresing et al. 2012

57 EUV waves Carley et al. 2013

58 Interplanetary CMEs

59 Outflows and magnetic clouds Spectroscopic measurements of outflow velocity have been used to compute more reliable estimates of magnetic flux in a CME Good agreement with estimates of flux in associated magnetic cloud. Harra et al., 2011

60 Models After eruption onset: Observations of flux ropes at 1AU, CME acceleration during flare impulsive phase, flare ribbons and flare arcade loops, EUV & X-ray dimmings. The above has led to the standard model Mandrini et al., 2005 Figure 1 from Shibata et al. (1995)

61 In 3-D Musset et al., 2015 Janvier et al Zharkov et al., 2011

62 Aly-Sturrock If all field lines simply linked to the solar surface, total energy of any force-free field cannot exceed that of an open field with the same flux distribution on the surface. Implies opening the field energetically unfeasible. Possible solutions: Reconnection Initial configuration not simply linked to the surface Initial field not force-free currents, gravity Only part of the field is opened 3D - flux rope can erupt by pushing the field aside.

63 Triggering and models Energy must come from the magnetic field Pre-CME equilibrium maintained by the balance of magnetic tension and magnetic pressure. Decrease of tension or increase of pressure could cause the pre-cme structure to seek a higher equilibrium and a CME may be triggered. Many different mechanisms proposed: Resistive Ideal

64 Resistive Models: Tether cutting Filament supported by field nearly aligned with the PIL Strongly sheared core field near the PIL Overlying less sheared arcade As shear increases reconnection starts. Tethers (AB and CD) are cut and AD and CB are formed. AD expand upwards and CB submerges, while sheared field near AD pulls the filament up. Overlying field is stretched, forming a current sheet. Doesn t address formation of the sheared field. A C D B Moore et al., 2001

65 Resistive models: flux cancellation By emergence: By reconnection: B.C. Low, 1996 van Ballegooijen & Martens, 1989 Debated! Currently favoured: accompanied by photospheric flux cancellation

66 How does the flux rope form? Successive reconnections along the PIL (van Ballegooijen & Martens, 1989) Is BPSS or HFT formed? Recent simulation of the process by Aulanier et al. (2010)

67 Resistive Models: Break-out Quadrupolar configuration with a null point above the central flux system. Shearing motions cause the central flux system to expand upwards, forming a current sheet at the null point. Reconnection starts, removing higher magnetic loops and allowing the core field to erupt. Kind of external tether cutting. Antiochos et al., 1999

68 Ideal models: Kink instability Can explain height-time profile of an erupting filament (Sakurai (1976)) Kink: critical twist 2π to 6π (Hood & Priest, 1979) Török & Kliem (2005) considered overlying field: Rapid decay with height leads to eruption Slow decay with height failed eruption

69 Fan & Gibson (2007)

70 Ideal models: Torus instability Current ring is unstable against expansion if external field decay is fast. Hoop force dominates over magnetic tension and flux rope can no longer be confined (Kliem & Török, 2005)

71 Fan & Gibson (2007)

72 Summary Flares and CMEs represent the largest release of magnetic energy within the solar system. Proximity of the Sun and many observing platforms allow detailed testing of ideas about how energy is released and transported. Knowledge can be used to inform understanding of other astrophysical environments. The standard model has many successful elements, and a 3-D picture is now well developed. Observational improvements show inconsistencies the problem isn t solved yet.

73