Parker Solar Probe and Solar Orbiter: Science Overview

Home , Granule (solar physics), Nanoflares

Marco Velli Earth, Planetary and and Space Sciences UCLA, JPL CIT [email protected]

Thanks: R. Lionello, L. Matteini, O. Panasenco, F. Pucci, F. Rappazzo, C. Shi, A. Tenerani, A. Verdini PSP Science Objec+ves Parker Solar Probe mission proﬁle

+ Venus Flyby #3 + Jul 6, 2020 + Venus Flyby #4 + First Perihelion + Feb 16, 2021 + at 35.7 RS + Nov 1, 2018 + Venus Flyby #7 + Nov 2, 2024

+ Launch + July 31, 2018 + Sun + Venus Flyby #5 + Oct 11, 2021 + Venus Flyby #6 + Mercury + Aug 16, 2023 + First Min Perihelion + at 9.86 RS + Venus Flyby #1 + Dec 19, 2024 + Sept 28, 2018 + Venus + Venus Flyby #2 + Dec 22, 2019 + Earth

At a speed of 190 km/s: probe crosses the corona taking about: 1 - 3 minutes s-granule PSP – SO Science Objec,ves

1. Trace the flow of energy that heats and accelerates the solar corona and solar wind. 2. Determine the structure and dynamics of the plasma and magne;c fields at the sources of the solar wind. 3. Explore mechanisms that accelerate and transport energe;c par;cles 1 What drives the solar wind and where does the coronal magne;c field originate from? 2 How do solar transients drive heliospheric variability? 3 How do solar erup;ons produce energe;c par;cle radia;on that fills the heliosphere? 4 How does the solar dynamo work and drive connec;ons between the Sun and the heliosphere? Focus today

Where does (Alfvénic) turbulence form and what is its role in coronal heating?

Is the wind source for fast and slow the same, and is there a steady component or is the wind always intermittent in nature?

Where does the heliospheric current sheet form and how stable is it close to Sun?

Reconnection and its role in coronal and inner heliospheric physics, MHD turbulence, nanoflares and coronal heating

Parker Solar Probe passes together with Solar Orbiter alignments, quadratures will help construct the history of plasma parcels from the corona into the heliosphere resolving origin evolution composition/ turbulence Solar corona, wind and magnetic activity Solar corona, wind and magnetic activity Sources of Energy

! c "! "! 1 S = E × B E = − V ph × B 4π c

Emerging Flux Waves and Turbulence 3 .WdmyrBh tal. et Wedemeyer-Böhm S. 338

Fig. 16 Schematic, simplified structure of the lower quiet Sun atmosphere (dimensions not to scale): The solid lines represent magnetic field lines that form the magnetic network in the lower layers and a large-scale (“canopy”) field above the internetwork regions, which “separates” the atmosphere in a canopy domain and a sub-canopy domain. The network is found in the lanes of the supergranulation, which is due to large-scale convective flows (large arrows at the bottom). Field lines with footpoints in the internetwork are plotted as thin dashed lines. The flows on smaller spatial scales (small arrows) produce the granulation at the bottom of the photosphere (z 0 km) and, in connection with convective overshooting, the weak-field “small-scale canopies”. Another result is the formation of the reversed granulation pattern in the middle= photosphere (red areas). The mostly weak field in the internetwork can emerge as small magnetic loops, even within a granule (point B). It furthermore partially connects to the magnetic field of the upper layers in a complex manner. Upward propagating and interacting shock waves (arches), which are excited in the layers below the classical temperature minimum, build up the “fluctosphere” in the internetwork sub-canopy domain. The red dot-dashed line marks a hypothetical surface, where sound and Alfvén are equal. The labels D-F indicate special situations of wave-canopy interaction, while location D is relevant for the generation of type-II spicules (see text for details). Please note that, in reality, the 3D magnetic field structure in the canopy and also in the subcanopy is certainly more complex and entangled than shown in this schematic sketch Temperature Profile

From Cranmer 2008 General Considerations on Turbulence in the Corona and Wind

(⇢U~ )= < ⇢U>~ r · r· < BB> B~ B~ B~ (⇢U~ U~ )= pT + + B~ + ⇢g < (⇢U~ U~ + ⇢... ) > r · r 8⇡ 4⇡ · r r· 4⇡ ✓ ◆

2 2 2 VT dU VT dA Vg R (U ) = U dr A dr 2 R2

@z± +(U V ) z± + z± (U V ) @t ⌥ a · r · r ⌥ a ± 1 1 T (z± z⌥) V U = P z⌥ z± <...>, ± 2 r · a ± 2 r · r Magne&c ﬁeld energy spectrum at 1 AU

Kiyani et al 2016 Turbulence (Helios results)

With well developed spectra, which evolve from a shape with P(ω ) ~ ω-1 at low frequencies, to P(ω ) ~ ω-1.67 at high frequencies.

Energy in the fluctuations E= ρu2/2+b2 /8π=E++E- also evolves with distance E(R) ~ R -a with a >3 - E+ ~ R -3.48 , E- ~ R -2.42

E- / E+ ~0.5 for R>2.5 AU Bavassano et al. 2000) Turbulence (models of the fluctuation amplitude) Accelera'ng Expanding Box Model

Tenerani&Velli, 17 Allows inhomegenous effects in MHD turbulence Complementary to moment models (Zank et al., Matthaeus et al., Tu and Marsch) Alfvén point provides dis0nct coronal turbulence regions Where%and%how%does%the%sharp%gradient%in%speeds% develop%close%to%the%Sun?%%

Antonucci et al., Telloni et al. 2007-2009

UVCS%%on%SOHO%suggest%slow/fast%separa6on%already%in%place%by%9%Rs% Where%and%how%does%the%sharp%gradient%in%speeds%develop% close%to%the%Sun?%%

Lionello et al. 2014 model shows the gradient as a function of distance from the sun : very similar. + Coronal hole boundary/S web “Blobs”(from(Heliospheric(current(sheet(( have(ﬂux9rope(like(nature( Periodic density structures

Structures with length scales of hundreds to several thousands of megameters and frequencies of tens to hundreds of minutes.

PDSs are formed in the solar corona as part of the slow solar wind release and/or acceleration processes. (Viall and Vourlidas 2015) Solar&Probe&will&directly&sample&wind&

Coronal Magnetic Topology: Streamers and Pseudostreamers. Fast, slow and hybrid wind. Presence of pseudostreamers (PS) a) extreme-high and extreme-low-proton-flux wind is associated with PS; "hybrid" type of outflow; b) intermediate proton speed but high electron temperature; c) spikes of proton density may represent PS plasma sheets; d) wind measurements; All above will allow to determine the structure and dynamic of the plasma and mag. fields at the sources of the solar wind. And also will answer the question how the processes in the corona affect the properties of the solar wind in heliosphere. Presence of streamers. a) not extreme-high densities and regular slow solar wind Current Sheet Stability in 2.5D in expanding box

x Region of Interest z (moving with the “blob”) y Wi Fast

F t = t ast 1 Wind

t = t2 > t1 B(x,y)

U0y(x)

Expanding Computational Box + Current sheets naturally arise in coronal magnetic fields

This configuration CAN NOT be imagined as a static equilibrium configuration. Almost ALL configurations must be intrinsically dynamic. Are there really FF fields or are all fields dynamic with low-level “turbulence”??? Parker, 1971 Syrovatskii 1978 + Heating the conﬁned corona

Reality

Rappazzo, Velli, Dahlburg Einaudi and permutations 2007-2016; Rappazzo & Parker 2013; Parker & Rappazzo,

Parker

REDUCED MHD SIMULATIONS Model + 3D Structure of Currents

+ J2/R Red = 1000 Yellow = 350 Max = 2.7x104 min = 0 + Parker & Rappazzo 2013

CURRENT SHEETS FORMATION IN TANGLED CORONAL MAGNETIC FIELDS

A. F. Rappazzo1,2 and E. N. Parker3 1 Bartol Research Institute, Department of Physics and Astronomy, University of Delaware, Newark, DE 19716, USA; [email protected] 2 Advanced Heliophysics, 1127 E Del Mar Blvd, Suite 425, Pasadena, CA 91106, USA; [email protected] 3 Enrico Fermi Institute, University of Chicago, Chicago, IL 60637, USA Received 2013 April 16; accepted 2013 June 24; published 2013 July 24

ABSTRACT We investigate the dynamical evolution of magnetic fields in closed regions of solar and stellar coronae. To under- stand under which conditions current sheets form, we examine dissipative and ideal reduced magnetohydrodynamic models in Cartesian geometry, where two magnetic field components are present: the strong guide field B0,extended along the axial direction, and the dynamical orthogonal field b.Magneticfieldlinesthreadthesystemalongtheaxial direction that spans the length L and are line-tied at the top and bottom plates. The magnetic field b initially has only large scales, with its gradient (current) length scale of the order of ℓb.Weidentifythemagneticintensitythreshold b/B0 ℓb/L.Forvaluesofb below this threshold, field-line tension inhibits the formation of current sheets, while above∼ the threshold they form quickly on fast ideal timescales. In the ideal case, above the magnetic threshold, we show that current sheets thickness decreases in time until it becomes smaller than the grid resolution, with the analyticity strip width δ decreasing at least exponentially, after which the simulations become underresolved. Key words: magnetohydrodynamics (MHD) – Sun: corona – Sun: magnetic topology Online-only material: animation, color figures

1. INTRODUCTION (MHD) turbulence, implicitly implying that current sheets thickness is limited only by numerical diffusion (i.e., resolu- All late-type main-sequence stars, for which the Sun is the tion) in dissipative MHD. But recent simulations of the decay prototype, emit X-rays (Gudel¨ 2004). Solar observations at of an initially braided magnetic configuration (Wilmot-Smith increasingly higher resolutions show that the X-ray corona has et al. 2009)haveshownthatinsomeinstancesthesystemforms structures at all resolved scales (Cirtain et al. 2013). only large-scale current layers of thickness much larger than Convective motions, which have more than enough energy to the resolution scale, in stark contrast with the recent result sup- heat the corona at temperatures >106 K, continuously shuffle porting the development of finite time singularities in the cold the coronal magnetic field line footpoints, giving rise to a plasma regime (Low 2013). magnetic field that is not in equilibrium (Parker 1972, 2000;van Furthermore, recent investigations suggest that the rate of Ballegooijen 1985). Parker (1972, 1988, 1994, 2012)pointedout magnetic reconnection can be very fast in low collisional that current sheets are an intrinsic part of the final equilibrium of plasmas, both in the MHD (Lazarian & Vishniac 1999;Lapenta almost all interlaced field line topologies. Thus the asymptotic 2008;Loureiroetal.2009;Huang&Bhattacharjee2010)and relaxation of the interlaced field to equilibrium necessarily the collisionless regime (Shay et al. 1999). Therefore, in order to involves the formation of current sheets, providing energy have an X-ray corona, it is critical that current sheets form, but dissipation presumably concentrated in small impulsive heating only above amagneticenergythreshold.Indeed,theenergyflux events, so-called nanoflares.Thispicturehashadastrong injected in the corona by photospheric motions is the average Poynting flux Sz B0 uph b /4π (Rappazzo et al. 2008, impact on the thermodynamical modeling of the closed corona ⟨ ⟩= ⟨ · ⟩ (Klimchuk 2006), but it is still controversial if and under which Section 3.1), that depends not only on the photospheric velocity circumstances current sheets form. uph and the axial guide field B0,butalsoonthedynamicmagnetic field b,andifdissipationkeepslowthevalueofb,theflux Sz Analytical models (van Ballegooijen 1985;Antiochos1987; ⟨ ⟩ Cowley et al. 1997)claimthatingeneralwell-behavedphoto- will be too low to sustain the corona (Withbroe & Noyes 1977). spheric motions will not lead to the formation of current sheets, In this Letter, we investigate, in a Cartesian model of the and that only a discontinuous velocity field can form disconti- closed corona, under which conditions current sheets form, and nuities in the coronal magnetic field, and counterexamples of their dynamical properties. well-behaved solutions of the magnetostatic equations have been reported (Rosner & Knobloch 1982;Bogoyavlenskij2000). 2. MODEL Alternatively, van Ballegooijen (1986)proposedthatthe random character of footpoint motions might generate, AclosedregionofthesolarcoronaismodeledinCartesian on timescales much longer than photospheric convection geometry as a plasma with uniform density ρ0 embedded in a strong and homogeneous axial magnetic field B0 B0 eˆz timescales, uniformly distributed small-scale current layers that = would heat the corona without forming discontinuous structures. well suited to be studied (e.g., see Dahlburg et al. 2012), as in Numerical simulations of boundary forced models (Einaudi previous work, with the equations of reduced MHD (RMHD). Introducing the velocity and magnetic field potentials ϕ and ψ, et al. 1996;Dmitruk&Gomez´ 1997;Rappazzoetal.2007)sug- 2 gest that the nonlinear dynamics of this system can be modeled for which u (ϕ eˆz), b (ψ eˆz),vorticityω ϕ, =∇× =∇× =−∇⊥ as a magnetically dominated instance of magnetohydrodynamic and the current density j 2ψ,thenondimensionalRMHD =−∇⊥ 1 + Sweet-Parker Current Sheets

+ Lv S = A ⌘

a a a 1/2 S L ⇠

Energy dissipated in an Alfvén time in the sheet is proportional to:

2 1/2 2 3 1/2 L vi⌧A = S L vAi⌧A = L S

“The observational and theoretical difficulties with the hypothesis of magnetic-field line annihilation suggest that other alternatives for the flare must be explored.” E. Parker, 1963 + The Plasmoid Instability (Tearing mode on Sweet-Parker)

+ L Tajima and Shibata 1997 Loureiro et al. 2007 Bhattacharjee et al. 2009 a Huang Y. M. and Bhattacharjee A. 2016 Huang et al. 2017

Renormalizing: 1/2 a 3/2 ⌧ S A ' L Sweet Parker ↵ a/L S ↵ =1/2 ⇠ ( 1+3↵)/2 ⌧ S A ⇠ ↵ =1/3 IDEAL TEARING + “Ideal” Tearing Mode Instability

+ • A critical aspect ratio (L/a)i must exist provides a “sup” for current sheets that can naturally form • Onset of “ideal” tearing can provide a scenario for the trigger of fast reconnection + The fractal reconnection scenario revisited +

Tenerani et al. 2015b, ApJ inspired by Shibata&Tanuma 2001 + The fractal reconnection scenario revisited + −0.000 0.02 τ t=18 A 0 0.01 −100 −0.001 −200 0.00

y/L −300 −0.002 −0.01 −400 −500 −0.02 −0.003 2.6 2.8 3.0 3.2 3.4 3.6 3.8

y 0.02 τ t=19 A 0 B −0.001 0.01 −0.004 −200 t=18τ 0.00 −400 A y/L

−0.003 −600 −0.005 t=19τA −0.01 t=19.8 −0.02 τA 2.6 2.8 3.0 3.2 3.4 3.6 3.8 −0.006 −0.005 0.02 τ −0.03 0.00 0.03 t=19.6 A 0 0.01 −0.007 −200

0.00 −400 −0.3 −0.2 −0.1 0.0 0.1 0.2 0.3 y/L y/L −0.01 −600 −0.02 Evidence of “recursive” tearing mode-like instabilities 2.8 3.0 3.2 3.4 3.6 0.02 τ during the nonlinear stage of a primary tearing mode t=20 A 0 0.01 −200 within a Harris current sheet. New plasmoids appear −400 0.00 −600 y/L to be generated, at each nth step, within smaller and −800 −0.01 −1000 smaller current sheets (CS), that consistently −0.02 j correspond to the inner layer of the (n-1)th unstable 2.9 3.0 3.1 3.2 3.3 3.4 3.5 z x/L CS. Tenerani et al. 2015b, ApJ inspired by Shibata&Tanuma 2001 + The fractal reconnection scenario revisited 13 1/3 S = 10 n =4 (1) an/Ln S ! ⇤ n 4 ⇠ 1/2 L1/L 6. 10 (2) an/Ln 1 Sn 1 ' ⇠ 6 1+(3/4)n (3/4)n L2/L 2. 10 Ln/L = S ,Sn = S ' ! 8 L /L 3. 10 3 ' (1) Threshold for instability growth 9 L4/L 10 (2) The nth unstable layer is the ' n 4 ⌧ ⌃ ⇤ ⌧ 10 ⌧ diffusion regions of layer (n-1)th tot ' 1 A,n ' A 1+(3/4)n n , ⌧ = ⌧ S ⌧ /S !1 A,n A ! A 4 Sn⇤ = 10 + Spectral Signatures of “recursive tearing”

−2 + 10 10−2 −3 10 10−3 Ec Em −4 10 1/3 10−4 −5 10 (1) an/Ln Sn 10−5 −6 ⇠ 10 y=0 10−6 y=0

−7 y=a/2 y=a/2 10 10−7 y=a y=a −8 + 10 y=2a 10−8 y=2a −9 10 10−9 1 10 100 1000 1 10 100 1000 k k

10−2 10−3 Ec Em 10−4 −2.3 k−1 k 10−5

10−6

10−7

10−8 10−9 + 1 10 100 1000 k Microstreams/plumes/jets

B C

Contribute to the fast wind? Conclusions