Category: UKSP Nugget

92. Null collapse and oscillatory reconnection about 3D magnetic null points

Author: Jonathan Thurgood (Northumbria University and University of Dundee), David Pontin (University of Dundee) and James McLaughlin (Northumbria University).

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Introduction

Magnetic fields play a key role in determining the dynamics of plasmas at all scales: from fusion experiments and laboratory plasmas to galaxies and accretion disks. For solar physics specifically, their importance cannot be overstated (19 past UKSP nuggests have ‘magnetic’ in the title or description!)  Much of our present understanding of reconnection has focused primarily on continuously driven (and hence at least quasi-steady) systems, while transient effects that are crucial in many plasma environments are typically neglected (though there are important exceptions). Furthermore, the study of 3D effects is still relatively uncharted territory. Combined with questions regarding the influence of chosen boundary conditions, the applicability of any given reconnection model to a particular solar case is often unclear. In our recent research, we have instead been studying the self-consistent formation and evolution of reconnection regions (current sheets) near different types of null points in response to finite energy perturbations in the form of MHD waves, isolated from the boundary effects. We typically find that the resulting evolution is characterised by phenomena which may be of importance for solar atmospheric dynamics. In this nugget, we present a case study which illustrates this typical behaviour.

Simulation Setup

We outline the dynamic evolution of localised and finitely-driven reconnection regions about fully-3D magnetic null points using numerical simulation of resistive, single-fluid MHD with the code Lare3d[1]. For the case study we consider a null of the form B = [x,y,-2z] (a so-called k=1 or azimuthal symmetric potential null). We subject the null to a flux-ring-like perturbation which corresponds to a (finite), nearly-uniform current distribution about the null which is localised in all three dimensions ([2] for details). The presumed physical origin of this perturbation is that of an accumulation of externally originating MHD waves impinging upon the null. This is a generic, well-studied phenomenon [3] that is known to evolve towards this state (i.e. our initial condition is a justified modelling simplification).

Initial Implosion (3D Null collapse)

The perturbation disrupts force-balance at the null and propagates inwards as a fast wave. It participates in an implosive process (null collapse) whereby energy is focused in increasingly small-scales towards the null. The precise behaviour of the implosion depends on the ability of the wave (its free energy) to continue to nonlinearly self- focus relative to plasma back-pressure and resistive diffusion, which both increase as plasma is increasingly compressed and current concentrations grow. Eventually, the process is halted due to a combination of these effects. The physics of this implosion is well-understood in the context of 1D harris sheets and 2D null points [4]. We have found that most aspects of 2D null collapse carry-over to various 3D nulls [5], and have derived empirical scaling relations and compared them to a reduced 1D analytical solution [4].

Here we choose parameters that correspond to a perturbation of sufficient energy to nonlinearly form a current sheet which is well-resolved on the grid at the time the collapse stalls. A representative current sheet evolution is shown in Figure 2, and the system of flow through the sheet at the halting time is shown in Figure 3. We stress that the sheet is localised in that it is detached from the influence of the boundary conditions and is not subject to continuous driving flows. The shock structure around the jets is a Petschek-like system – replete with a collimated jet bound by standing slow shocks and fast termination shocks – as confirmed by detailed analysis of the jump conditions in [4].

Reconnection Reversals

As the outflow evolves, we observe the pooling of hot over-dense plasma at the heads of the reconnection jets, leading to back-pressures that choke off the outflow. This is accompanied by the relief of compression in the current sheet itself via the expulsion of the excess plasma. A plasma rarefaction forms in the inflow lobes, with plasma accelerated through the current sheet unreplenished in the absence of continual driving flows. The consequence of this plasma and flux redistribution is that the balance of forces acting across the spine and fan planes (the collapse of which is necessary to sustain the reconnection and current sheet) must change as reconnection proceeds. The competition of these effects sets up secondary, oppositely-directed null point collapses which leads to successive changes of the current sheet orientation and polarity (‘reconnection reversals’ or ‘current sheet reversals’). Naturally, the concomitant reconnection also reverses its directionality – thus we have time-dependent, periodically-reversing reconnection, i.e. 3D oscillatory reconnection (OR) . We find that in 3D OR, the magnetic reconnection is of the spine-fan type (multiple types of reconnection are permitted in 3D as per [6]) as visualised in Figure 4.

Wave Output

In these simulations the reconnection only occurs in a small region near the null point. However, the global effects of the magnetic field restructuring are not artificially confined to this area by computational boundaries, but rather may escape the vicinity of the current layer as freely-propagating waves. We find many escaping MHD waves are generated and highlight a selection in Figure 5. The time-distance diagrams for velocity and density signal along the spine axis are consistent with the intermittent launching of Alfvén (A) and Slow (S) waves at each reversal. The Alfvén waves manifest as propagating, counter-rotational vorticity tubes, and slow waves are large-amplitude, anharmonic pulses which propagate mass away from the reconnection jets along the spine and fan field lines (see [2]). Fast wave pulses (potentially shocks) are also ejected from the edge of the current sheet after each successive reversal (see [4]).

Summary

The above case-study is typical of our findings, where the evolution is characterised by: (1) reconnection that is time-dependent and periodic (exhibits oscillatory reconnection) and (2) the periodic generation of propagating MHD waves in a variety of modes that escape from the vicinity of the reconnection site. The direct production of observable periodic phenomena by inherently time-dependent reconnection is a tantalising prospect for unexplained periodicities in solar physics that are associated with reconnection events, such as QPPs in flares (see [7]).… continue to the full article

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91. Magnetic field line tangling and twisting in the corona

Author: Simon Candelaresi, David Pontin, Gunnar Hornig at the University of Dundee, Anthony Yeates at Durham University, Paul Bushby at Newcastle University.

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Plasma motions on the Sun’s photosphere can lead to tangling of the coronal magnetic field lines that are rooted there. The high tangling can generate thin current layers, potential sites of magnetic field reconnection and energy conversion. Here we study the field line tangling induced by these photospheric motions.

Introduction

The solar surface (photosphere) is a highly dynamical and turbulent plasma that exhibits turbulent flows. These are due to convective motions in the outer solar interior that transport heat to the surface. On the surface, this plasma cools through thermal radiation and sinks back where it is reheated.

Further up in the solar atmosphere we observe coronal magnetic fields that are anchored at the photosphere. The presence of turbulent photospheric flows means that the field lines in the corona are continually reorganised, twisted and even braided. Such changes in the magnetic field’s topology have implications for energy transport from the photosphere, the formation of thin current layers in the corona, and subsequent dynamic phenomena involving magnetic reconnection.

Scenarios

From observations we know the time evolution of the line of sight magnetic field on the photosphere. Using the Fourier Correlation Technique [1, 2] we can extract the horizontal component of the plasma velocity on the photosphere. Of course, observations of the solar magnetic field are strongly limited in resolution. Furthermore, the inferred horizontal motions strongly depend on the method used for extracting the horizontal velocities. Therefore, we also consider here velocities obtained from magneto-convection simulations [3, 4], which are known ‘exactly’. Finally, we compare with a well known benchmark motion that leads to efficient fluid mixing, called the blinking vortex. This consists of two overlapping regions along the x-axis where we apply alternating twisting motions of opposite signs (for details see [5, 6]).

As a proxy for the tangling of field lines in the corona, we can measure the mixing of fluid particles by each of our model photospheric flows by following ‘marked’ particles. We mark the particles by using an initial passive scalar distribution of the form c(x, y) = x + y. As the particles move, this distribution will change. The mixing is visualised for the flows from the magneto-convection simulations in Figure 1. There we plot the passive scalar distribution, as the motions turbulently mix the plasma. There is clearly a strong mixing that we will quantify using the winding number and the topological entropy.

Winding Number

If we want to know the efficiency of the winding of the surface motions we need to integrate the velocities of test particles in time. With that we calculate for every flow trajectory the average winding of other trajectories around it which gives us an averaged winding number (here we imagine a domain with three dimensions: the two flow dimensions plus time). Doing this we find mean winding numbers of 0.122 for the magneto-convection simulation, 0.05 for the photospheric observations and 0.0243 for the benchmark blinking vortex. For the observed flows on the real photosphere, this means that the average winding increases by an angle of 0.63 degrees per hour. These calculation involve a normalisation in order to obtain dimensionless numbers. For details refer to [6]. Although the used benchmark exhibits a highly efficient mixing, the simulated and observed photospheric motions are more efficient in mixing the fluid.

Topological Entropy

Another way of quantifying the tangling of a 2d fluid or 3d magnetic field lines, is the finite time topological entropy [7-9]. It measures the exponential stretching rate of a material line that is being advected by the fluid. An example of such an advected material line is shown in Figure 2. Similar to the experiments using the winding number, the simulated convection layer and the observations show a high degree of tangling in terms of finite time topological entropy with values of 2.078 for the magneto-convection simulations, 1.598 for the observed photospheric flows and 0.4928 for the benchmark flow (see Figure 3). A Value of 1 means that in the course of one (normalized) time unit (4.696 hours for the observations) a material line is stretched by a factor of e.

Conclusions

The observed photosphere and simulated convection zone are highly efficient in mixing the photospheric plasma – even more so than the benchmark flow, itself known to be a very effective mixing flow. For anchored magnetic flux tubes this means that on the time scale of approximately 3 hours we can expect a similar complexity as seen in the so called magnetic pig-tail braid that consists of alternating positive and negative twisting regions. With respect to the Parker braiding problem our result suggest that the present motions have the potential to lead to highly tangled and braided magnetic fields.

References

  • [1] Yeates, A. R., Hornig, G., & Welsch, B. T. 2012, A&A, 539, A1
  • [2] Welsch, B. T. 2015, Publ. Astron. Soc. Jpn., 67, 1
  • [3] Bushby, P. J., Favier, B., Proctor, M., & Weiss, N. 2012, Geophys. Astrophys. Fluid, 106, 508
  • [4] Bushby, P. J., & Favier, B. 2014, A&A, 562, A72
  • [5] Aref, H. 1984, J. Fluid. Mech., 143, 1
  • [6] Candelaresi, S., Pontin, D. I., Yeates, A. R., Bushby, P. J., & Hornig, G. 2018, arXiv:1805.030101
  • [7] Adler, R. L., Konheim, A. G., & McAndrew, M. H. 1965, Trans. Amer. Math. Soc., 114, 309
  • [8] Mangalam, A., & Prasad, A. 2017, Adv. Space Res.
  • [9] Candelaresi, S., Pontin, D. I., & Hornig, G. 2017, Chaos, 27, 093102

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90. A novel magneto-seismology technique for solar magnetic field diagnostics

Author: Matthew Allcock and Robertus Erdélyi at the University of Sheffield.

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Introduction

The magnetic field in the solar transition region and corona is amongst the least understood but the most significant features in solar physics. As the coronal magnetic field is relatively weak and the plasma is optically thin, it is difficult, and often impossible, to directly measure the magnetic field vector field in the transition region and corona using traditional methods such as Zeeman splitting [1]. Thankfully, we can harness proxy information to do a better job by means of solar magneto-seismology (SMS).

Magnetohydrodynamic (MHD) waves, that permeate almost all solar structures, provide us with one such proxy. Ingenious SMS techniques have utilised observed MHD wave parameters including frequency, period ratio, and damping rates to diagnose the magnetic field in a variety of solar atmospheric  structures. One characteristic of MHD waves in magnetic waveguides that has not yet been utilised is how the associated wave power is distributed across the structure (the transverse eigenfunction) – in particular, the (a)symmetry of the waves.

What are asymmetric MHD waves?

Theoreticians like symmetry – it makes their lives easier. As a consequence, we have become familiar with wave mode identification in terms of the traditional kink and sausage MHD modes in idealised and symmetric structures. For example, see Figure 1, show a symmetric plasma slab oscillating with a symmetric kink and sausage surface modes, respectively. These animations were created using Mayavi2 and over 100 more animations of MHD eigenmodes can be found at the SP2RC research group website.

These wave modes map reasonably well to observations, but are based on the simplistic assumption that the waveguide is symmetric (reflectional symmetry for Cartesian structures and axisymmetry for cylindrical structures). It is likely that the highly inhomogeneous solar atmosphere supports structures that are asymmetric, such that the plasma parameters on one side of the structure differ from those of the other. Likely candidate structures are elongated magnetic bright points (see Figure 2), prominences, and light walls, which are ensembles of surges and other magnetic phenomena above sunspot light bridges (see Figure 3). MHD waves guided by these asymmetric waveguides can therefore also be asymmetric.

We categorise these asymmetric MHD waves as either quasi-kink or quasi-sausage modes, depending on whether the opposite boundary oscillations are in phase or in anti-phase [5]. This nomenclature is chosen to emphasize the difference to their symmetric counterparts, namely that several properties intrinsic to symmetric kink and sausage modes are blended when this symmetry is broken. For example, the quasi-kink mode does not conserve the cross-sectional width of the waveguide.

Asymmetric eigenmodes can exist as either surface or body modes, depending on the equilibrium parameter regime of the waveguide, its environment, and the form of the initial perturbation [5]. For example, shown in Figure 5, the first two orders of fast quasi-kink and quasi-sausage body modes, where we also depict areas of high and low perturbed density in red and blue. Note that body modes are not significantly affected by asymmetry in the background plasma.

A novel technique for solar magneto-seismology

Varying the strength of the magnetic field within an asymmetric waveguide changes the degree of asymmetry of the wave. In other words, asymmetry of an observed wave contains information about the underlying magnetic field. Therefore, if we can measure this asymmetry, we can approximate the magnetic field by solving an inverse problem.



Figure 5. Quasi-kink (left column) and quasi-sausage (right column) body modes of an isolated slab of plasma. The areas of high and low perturbed density in red and blue. The red tubes depict the internal magnetic field lines. Click for the movies.

How can we quantify the asymmetry of MHD waves? One method that we have developed at SP2RC involves taking the ratio of the amplitudes of the transverse oscillations on the structure’s opposite boundaries.

We model Cartesian structures in the solar atmosphere by a magnetic slab of plasma embedded in an asymmetric background plasma . By taking the amplitude of transverse oscillation, which is distributed non-uniformly across the structure, one can find that the amplitude ratio as a transcendental function of the density, temperature, frequency, wavenumber, and internal magnetic field strength, whose explicit form is given in [6], and depends on the type of wave mode observed.

In theory, the amplitude ratio, period, and wavelength of the asymmetric wave, and the density, temperature, and spatial scale of the background structure, can all be measured (granted, some more easily than others). From the amplitude ratio, one may make an estimate of the magnetic field strength within the structure. For wavelengths much longer than the width of the structure, this inversion can be completed analytically to give the magnetic field strength inside the structure. Alternatively, and more accurately, the inversion can be made without approximation using a simple numerical procedure [6].

Conclusions

We have developed a technique that, for the first time, exploits observations of asymmetric MHD waves to diagnose the solar atmosphere’s elusive magnetic field. We anticipate that the improved spatial resolution of next-generation solar instrumentation, including DKIST and EST, will be sufficient to resolve the cross-sectional wave distribution of solar structures sufficiently to employ the amplitude ratio technique described above.

The main limitations with this technique lie in its simplicity – alas, solar atmospheric structures are not so simple! The infancy of this technique’s development mean that there are a number of exciting avenues for improvement, making it a promising new technique for furthering our understanding of solar magnetic fields.

Acknowledgements

M. Allcock acknowledges the University Prize Scholarship from the University of Sheffield. R. Erdélyi acknowledges the support from the UK Science and Technology Facilities Council (STFC, grant number ST/M000826/1) and the Royal Society.

References

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89. When are Coronal Jets miniature CMEs?

Author: Peter Wyper at Durham University, with C. Richard DeVore and Spiro Antiochos at NASA GSFC.

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Introduction

Coronal jets have been the subject of study for decades, particularly following their discovery in X-ray images in the 1990s [1]. They consist of a brightening of small coronal loops at their base, and a jet of heated plasma launched upward along open magnetic field lines. A burst of magnetic reconnection between the short base loops and the surrounding field is the obvious candidate for explaining both features [1], and many jets fit this simple picture. However, a sub-set of jets have a more complex evolution; involving the ejection of cool, dense plasma with clear untwisting motions. A simple burst of reconnection cannot easily explain them [2]. Something more complex is required.

Relatively recent observations have revealed that this sub-set of jets usually involve the eruption of a so called “mini-filament”, a small-scale version of their larger cousins which erupt to produce Coronal Mass Ejections (CMEs) [3]. A comparison of a mini-filament jet with a full-scale filament eruption is shown in Fig. 1. The similarities between the tiny and large-scale filament eruptions suggests the two phenomena may be related, i.e. that these jets are miniature CMEs.

Recently, we developed a numerical model for mini-filament jets and showed that they can indeed be thought of as scaled down CME eruptions [5]. Thus, coronal jets can be thought of as miniature CMEs when they are generated by the eruption of a mini-filament. In this nugget, we discuss this model and how different realisations of the model can help explain the different jet dynamics in the run up to these tiny eruptions.

Scaling down the eruption process

Our model is essentially a scaled down version of the Breakout Model for large scale CMEs (see [6] and [7] for a full description). Figure 2 depicts schematically the different stages of the model. The initial magnetic field (Fig. 2a) consists of a 3D null point (zero point in the field) with a domed separatrix surface (separating open and closed field) and inner and outer spine lines. Starting with this field there are then 3 main stages:

  1. Filament channel formation: Some process forms sheared field lines (yellow) capable of supporting mini-filament material in the closed field. In our numerical model we use surface motions, but in principle this could be flux emergence or shearing and cancellation. The system remains in force balance since the magnetic tension of the cyan field lines (the strapping field) balances the outward magnetic pressure of the sheared field lines (the filament channel).
  2. Breakout: The increasing magnetic pressure of the growing filament channel expands the strapping cyan field lines. Eventually, this forms a current layer at the null above, which begins to reconnect the strapping field. This reduces the downward force on the filament and allows it to expand faster. Beyond a tipping point, the feedback between the reconnection at the null and the upward expansion of the filament channel becomes self-sustaining and the filament starts to rise. This is the so called Magnetic Breakout mechanism [6]. As the filament channel stretches upwards at a steady pace, a second current layer forms beneath it. Additional reconnection in this layer turns the sheared field lines into a twisted flux rope.
  3. Eruptive jet: Finally, when the rising flux rope reaches the null current layer, it is reconnected on to open field lines. The transfer of twist on to open field drives non-linear Alfvén waves that create an untwisting jet.

This multi-step model explains how the mini-filaments erupt, how they produce jets and why the jets have untwisting motions. It also theoretically links these tiny eruptions with large-scale filament eruptions through the same eruption mechanism (for a detailed discussion see [5], [8]).

Explaining different observed behaviours

Are the different stages of the model observed though? The model places no constraints on where the mini-filament comes from and agrees well with the ultimate untwisting jet. The interesting thing observationally is whether reconnection outflows are observed before the main untwisting jet that match those we expect in the breakout phase. Some mini-filament jets do show this, particularly when the surrounding field is highly inclined [9]. However, others in less inclined field regions do not. This suggests that the background field inclination is an important factor.

To better understand this, we used the Adaptively Refined Magnetohydrodynamics Solver (ARMS) code to run three simulations with different background field inclination angles: +22, 0 and -22 degrees. The angle is with respect to the vertical direction, Fig. 2a. Figure 3 shows the velocity magnitude during the jet eruption with vertical field. After a short burst a long, slow breakout phase occurs with weak vertical exhaust flows. For the other two cases we find faster, denser reconnection exhausts over a shorter period. The main reason is that in the non-vertical cases the angle of contact between the strapping (cyan) and overlying (red) field lines is increased, driving more intense reconnection at the null (for full details see [8]). Thus, for higher inclination angles what are likely to be more observable reconnection exhausts are produced, agreeing qualitatively with the observations.

Conclusions

Coronal jets are deceptively simple events that we are only now getting a clear understanding of. It is now apparent that in the cases where a mini-filament is involved they share similarities with large-scale filament eruptions. Such jets can therefore be thought of as mini-CMEs. We have introduced a model for jets driven by mini-filament eruptions that explains the link between these large-scale and small-scale events. Our work suggests that eruptions across vastly different scales in the corona can be understood within the same framework. The three realisations of the model we studied also help explain the differences between individual mini-filament jets. With further space missions on the horizon, particularly the improved views of the poles with Solar Orbiter, the future looks bright for further unpicking the secrets of these small but mighty events.

References

  • [1] Shibata, K.

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88. Excitation of coronal loop oscillations by coronal rain

Author: Petra Kohutova and Erwin Verwichte at the University of Warwick.

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A little bit of background

Coronal rain is a common phenomenon occurring in active region coronal loops. As the name suggests, it shares many parallels with its terrestrial counterpart; it consists of cool plasma condensations falling from coronal heights towards the solar surface guided by magnetic field lines. Coronal rain is formed as a direct consequence of thermal instability. A coronal loop is likely to become thermally unstable if it is heated predominantly at the footpoints. If the thermal conduction along the loop is not efficient enough, the radiative losses from the plasma at the loop top can overcome the heating input from the footpoints, resulting in onset of a thermally unstable regime and in the runway cooling of the plasma at the loop top. This leads to the formation of cool and dense plasma condensations, which then fall towards the solar surface in the form of coronal rain showers. Recent high resolution observations have shown that coronal rain is much more common than previously thought [1], suggesting it has an important role in the chromosphere-corona mass cycle. Due to its origin, coronal rain also provides us with physical insight into the atmospheric thermal cycle.

What have we observed?

We analysed observations of a thermally unstable loop taken by IRIS (Fig. 1). The loop shows a significant amount of coronal rain formation. As soon as the coronal rain appears, the loop starts to oscillate in the direction parallel to the plane of the loop. From time-distance plots we can deduce that the coronal rain formation excites a vertically polarised kink mode (Fig. 2). This is not surprising, as abrupt formation of a cool and dense condensation region at the loop top pushes the loop axis downwards. Lorentz force acts to restore the loop shape and the loop starts to oscillate as a result. The lack of observable damping suggests that the excitation is prolonged in time. We further observed a decrease in the loop oscillation period. This is a result of the evacuation of the loop top as the coronal rain drains mass from the loop that then falls towards the footpoints. If the mass of the loop changes, we expect a shift in the period of the fundamental kink mode. Based on the observed change of the period, we can estimate the total mass of the plasma lost due to coronal rain to be a third of the total loop mass (see [2] for a complete discussion). We also observed a change in the oscillation amplitude. Using a relation between oscillation amplitude and the mass of the condensation plasma derived in [3], we arrive at total rain mass consistent with the estimate based on the change in period.

Can we simulate this?

Yes! We carried out 2.5D MHD simulations of the above scenario using Lare2d code. The problem is set up to include a realistic stratified atmosphere consisting of a cool chromosphere, a thin transition region layer and a hot corona. We used current-free potential magnetic field configuration representative of a coronal arcade and created a long and thin loop that is denser than the surrounding plasma. Finally, we introduced a cool and dense condensation region at the loop top (see [4] for a complete description of the setup). As in the observations, the presence of the condensation region is found to excite fundamental harmonic of a vertically polarised kink mode in the loop. We also managed to reproduce the decrease in the oscillation period of the loop, which can be seen in the simulation as the condensations leave the loop top and fall into loop legs (Fig. 3). Considering that the simulation is set up such that the total mass of the plasma in the loop is conserved, the period shift in this case is caused by the redistribution of the mass along the loop leading to the change in the longitudinal density profile.

Why is this important?

This mechanism can explain excitation of vertically polarised loop oscillations without the need for placing specific external drivers below the loop. Continuous rain formation leading to prolonged excitation can explain the apparent absence of damping. Our work therefore shows that we need to be careful when drawing conclusions from the observations of “decayless” small-amplitude coronal loop oscillations commonly seen in AIA [5,6] and linking them to photospheric modes, as they can be excited by thermal instability in the corona. Finally, we showed that the oscillations excited by coronal rain have seismological potential. The evolution of the oscillation period can be used to determine the fraction of the loop plasma mass that becomes unstable. This gives us information about the parameters of the heating function and hence tells us how localised the heating of the coronal loop is.

References

  • [1] Antolin, P. & Rouppe van der Voort, L. 2012, ApJ, 745, 152
  • [2] Verwichte, E. & Kohutova, P. 2017, A&A, 601 L2
  • [3] Verwichte, E., Antolin, P., Rowlands, G., Kohutova, P., & Neukirch, T. 2017, A&A, 598, A57
  • [4] Kohutova, P. & Verwichte, E. 2017, A&A, 606 A120
  • [5] Wang, T., Ofman, L., Davila, J. M., & Su, Y. 2012, ApJ, 751, L27
  • [6] Nisticò, G., Nakariakov, V. M., & Verwichte, E. 2013, A&A, 552, A57

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87. Giant solar loops and LOFAR radio observations

Author: Hamish Reid and Eduard Kontar
Astronomy and Astrophysics Group, University of Glasgow.

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Introduction

Solar flares accelerate particles that then produce radio emission as they move through the corona. Type III radio bursts are commonly observed signatures of electron beams propagating along magnetic field that extends into interplanetary space. Seen less often are radio U-bursts and J-bursts, signatures of electron beams propagating along magnetic loops confined to the corona [1,2]. It is not clear why we do not commonly see U-bursts or J-bursts. To answer this question, we need to understand the properties of the accelerated electrons and the magnetic loops they travel along.

Observations

We used Low Frequency Array (LOFAR) imaging spectroscopy [3] to analyse two solar radio J-bursts and one U-burst from a storm of bursts that occurred shortly after a large flare at 12:00 UT on 6 May 2015. The dynamic spectrum of all three bursts is shown in Figure 1. The inverted U or J shape arises from an electron beam travelling up the ascending leg of a magnetic loop and through a decreasing plasma density, corresponding to a decrease in the radio frequency of the burst with time. But unlike the more common Type III bursts, in the U and J-bursts the electrons then travel back down the descending leg of the magnetic loop, through an increasing plasma density, corresponding to an increase in the radio frequency of the burst with time. An inverted J shape is made instead of an inverted U shape if the beam stops emitting radio waves at the apex of the loop.

Low Frequency Array (LOFAR) observations have a substantially improved frequency coverage for imaging than past observations, combined with sub-second time resolution. We do not see large numbers of J- and U-bursts, especially at high frequencies. What is new with LOFAR is, for the first time, to be able to REALLY observe (image) the magnetic loop, with many frequencies. Below 100 MHz we previously had a maximum of three frequencies, which is not great for deducing plasma properties.

Figure 2 shows an image of the U-burst between 75 and 40 MHz at the times of peak intensity, so the higher frequency contours were earlier in time than the lower frequency ones. The bursts allow a large part of the magnetic loop to be visible at altitudes not dense enough for EUV or X-ray imaging. This U-burst showed faint radio emission originating from the descending leg of the magnetic loop. A fit to the radio centroids finds a loop with an altitude of approximately 1 solar radius and a length around 1.5 solar radii from the bottom to the apex of the loop. Starting heights for the radio emission were between 0.6 – 0.8 solar radii. The magnetic loop model was combined with the frequency evolution in time to estimate exciter velocities between 0.13c and 0.23c, all without requiring the common assumption of a coronal density model or emission mechanism. We also estimated a density model of the magnetic loop. It had a much smaller density gradient than the standard density models of the quiet Sun, although there are uncertainties in the radio positions because of scattering effects of the light from the source to the observer [4].

Occurrence of U-bursts

Why are U-bursts and J-bursts not observed more often than their type III counterparts? For an electron beam to generate radio emission, it must become unstable and generate Langmuir waves during propagation. The Langmuir wave instability time must be shorter than the propagation time or no Langmuir waves will be generated. We used the results above to construct a graph of instability time versus propagation time, shown in Figure 3.

As the electron beam travels through the solar atmosphere, the bump-in-tail instability can cause Langmuir waves to be generated. This requires an instability distance [5,6] which is independent of initial beam density. If the magnetic loop is too small then the electron beam will not become unstable and no radio emission will be generated.

Once the electron beam is unstable to Langmuir waves, the instability time must be short. This timescale is proportional to the square root of the background electron density and inversely proportional to the electron beam density [7,8]. So higher initial beam densities correspond to shorter instability times (i.e. smaller timescales). However, the low magnitude of the density gradient in closed loops (J-, U- bursts) keeps the background density high and causes instability times to be longer than for open loops (type III bursts).

Conclusions

The fine spectral and temporal resolution of the LOFAR images between 75 and 40 MHz indicated a loop-shaped structure extending from the flaring active region in which an X-ray source was present [3]. Using this enhanced resolution, we extracted properties of both the accelerated electron population and the background plasma they travelled through.

We found that U-bursts or J-bursts are only produced from a restricted range of accelerated beam and background plasma parameters, resulting in type III bursts being more frequently observed. The large instability distances required before Langmuir waves are produced by some electron beams, and the small magnitude of the background density gradients, makes closed loops less fertile for producing radio emission than loops that extend into interplanetary space.

References

  • [1] Maxwell, A., & Swarup, G. 1958, Nature, 181, 36
  • [2] Fokker, A. D. 1970, Sol. Phys., 11, 92
  • [3] Reid, H. A. S., & Kontar, E. P. 2017, A&A, 606, A141
  • [4] Kontar, E. P., Yu, S., Kuznetsov, A. A., et al. 2017, Nature Comms, 8, 1515
  • [5] Reid, H. A. S., Vilmer, N., & Kontar, E. P. 2011, A&A, 529, A66
  • [6] Reid, H. A. S., Vilmer, N., & Kontar, E. P. 2014, A&A, 567, A85
  • [7] Vedenov, A. A. 1963, Journal of Nuclear Energy, 5, 169
  • [8] Kontar, E. P. 2001a, Sol. Phys., 202, 131

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86. Evidence of recurrent reconnection driving fan-shaped jets.

Author: Aaron Reid, Mihalis Mathioudakis (Queen’s University Belfast), Vasco Henriques (UiO, Norway), Tanmoy Samanta (Peking University, China).

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Introduction

On-disk fan-shaped jets most often appear at light bridges ([1],[2],[3]), where the dominant magnetic field orientation changes from vertical to horizontal [4]. This has led to the theory that fan-shaped jets are formed via shearing reconnection between the vertical and horizontal photospheric magnetic field lines at the light bridge – umbra interface [5].

Three-dimensional simulations have reinforced this theory, showing fan-shaped jets being formed during shearing reconnection between horizontal magnetic fields and a vertical current sheet, with the base of the jets corresponding to the location of magnetic reconnection [6]. While the initial acceleration of the jets is due to the magnetic tension, the kinematics associated with this phenomenon are thought to be due to the gas pressure gradient created by the heating at the base of the jet [7].

In this nugget, we present evidence of recurrent magnetic reconnection driven by running penumbral waves in the photosphere. This reconnection at the penumbral edge leads to the appearance of chromospheric fan-shaped jets.

Limb observations of fan-shaped jets

We acquired Hα imaging spectroscopy data of a sunspot near the solar limb using the Swedish 1-m Solar Telescope (SST). At the sunspot edge, fan-shaped jets can be seen in the Hα line core (top-middle panel of Fig. 1). The Hα red wing shows the descent of previous jets, while in the blue wing, these events are not well captured. This is most likely due to line-of-sight effects Doppler-shifting the falling material. Using the SDO alignment tool of Rob Rutten, co-spatial and co-temporal SDO data was also created, and are also shown in Fig. 1.

The red box in Fig. 1 highlights a region where repetitive brightenings exist in the wings of Hα, indicating photospheric temperature enhancements. These brightenings have line profiles matching that of nearby Ellerman Bombs, which are caused via photospheric reconnection [8]. This agrees with the current theory of fan-shaped jet formation [6]. Light-curves were created for the acquired datasets, and it was apparent that these intensity enhancements at the base of the fan-shaped jets only occur in the 1600A and 1700A channels of AIA and the Hα line wings, and so are solely photospheric. The average time between brightenings was 214 seconds, which is similar to typical periods of running penumbral waves [9]. The peaks in the photospheric light-curves are also near-simultaneous to the appearance of new fan-shaped jets in the chromosphere. Fig. 2 shows how the photospheric brightening directly relates to the fan-shaped jets in the chromosphere. The bottom panels show how the brightening erupts material into the chromosphere, which is co-spatial and co-temporal to the appearance of a new bright-jet front in the Hα line core.

Curvilinear time slices were created for the fan-shaped jets, as highlighted by the green line in Fig. 1. The resultant output is shown in Fig. 3, and the bright front is clearly visible in the Hα line core, along with SDO AIA 304A, 171A, and 131A. Interestingly, the body of the jets appear dark in all AIA channels bar 304A (and 1600/1700A, but the jets are not seen in these channels). This darkening is most likely optically thick plasma caused by strong absorption due to an increased density of hydrogen and singly ionised helium [10].

We tried to characterise the fan-shaped jets by applying polynomial fitting to the jet front detected in the Hα red wing. This would enable a velocity and deceleration profile to be obtained for the falling material. The resultant fit is shown by the blue dashed line in Fig. 3, where the jet material appears to be accelerating back to the solar surface at a rate consistent with solar gravity. However, the initial rising of the jet fronts is clearly not decelerating in the same way. Rather, the jet fronts appear to protrude at a constant velocity of 30 km/s according to linear fitting attempted on the brightest pixels of the jet front in Hα. This suggests the presence of a driver counteracting the effects of solar gravity.

Conclusions

We have shown evidence for recurrent photospheric reconnection leading to the appearance of fan-shaped chromospheric jets at a sunspot edge. While fan-shaped jets are most commonly observed across light-bridges, the current theory for their appearance allows them to form anywhere strong vertical and horizontal fields may interact. In this case, the strong vertical fields from the umbra interact with the nearby horizontal magnetic fields, causing reconnection at the edge of the penumbra.

The fan-shaped jets appear with a bright front in the low temperature chromospheric channels, while in the hotter coronal channels, darkness is left behind in the body of the jets. This implies a relatively hot jet front (~50,000K) with cool, optically thick plasma left in the wake. Our analysis of the jet protrusion implies that there is a driver to the emanation of the jet. The Fourier power spectra of the sunspot penumbra hint that the peak frequency corresponds to the average time between photospheric eruptions at the base of the jets, implying that running penumbral waves could be such a driver.

References

  • [1] Asai, A., Ishii, T. T., & Kurokawa, H. 2001, ApJ, 555, L65
  • [2] Robustini, C., Leenaarts, J., de la Cruz Rodriguez, J., Rouppe van der Voort, 2016, A&A, 590, A57
  • [3] Roy, J.-R. 1973, Sol. Phys., 32, 13
  • [4] Shimizu, T., Katsukawa, Y., Kubo, M., et al. 2009, ApJ, 696, L66
  • [5] Robustini, C., Leenaarts, J., & de la Cruz Rodriguez, J. 2017, arXiv:1709.03864
  • [6] Jiang, R.-L., Shibata, K., Isobe, H., Fang, C. 2011, Research in Astronomy & Astrophysics, 11, 701
  • [7] Li, Z., Fang, C., Guo, Y., et al. 2016, ApJ, 826, 217
  • [8] Reid, A. L., Mathioudakis, M., Doyle, J. G., et al. 2016, ApJ, 823, 110
  • [9] Lohner-Bottcher, J., & Bello Gonzalez, N. 2015, A&A, 580, A53
  • [10] Anzer, U., & Heinzel, P. 2005, ApJ, 622, 714

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85. The role of the magnetic field in sunquakes

Author: Lucie Green, Gherardo Valori, Francesco Zuccarello, Sarah Matthews at MSSL UCL, Sergei Zharkov at University of Hull and Salvo Guglielmino at Università degli Studi di Catania.

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Introduction

Sunquakes are sporadic bursts of sound waves that are generated at the surface of the Sun as a side-effect of solar flares and coronal mass ejections. Sunquakes are thousands of times more energetic than the earthquakes seen on our planet and must be powered by a huge energy reservoir which is thought to come from the magnetic field in the solar atmosphere. But there is a mystery around how the energy released during a solar flare and a CME is transported from the atmosphere into the Sun and onto the epicentre that is the origin of the sunquake. While energy transport has been proposed to occur via particles, waves or a bulk motion of the magnetic field, no real explanation is available as to why sunquakes occur exactly where they do.

With the launch of SDO we gained the opportunity to study the evolution of the corona in great detail in the time leading up to and during sunquakes, and relate this evolution to the generation of sunquakes. In this nugget we show the results of our investigation of the 3D magnetic field in the region of a sunquake that occurred on 17th February 2013. We show that a special configuration of the magnetic field may have focused energy onto a particular location in the lower atmosphere, thereby producing the sunquake [1].

Events in the solar atmosphere

The sunquake of 17th February 2013 occurred in NOAA active region 11675. The evolution of this active region in the time leading up to the sunquake is dominated by the emergence of a series of magnetic bipoles with their associated motions, collision of polarities and, in one bipole, the splitting of the positive polarity. Collectively, this evolution creates a complex magnetic field configuration with regions of field that are highly sheared.

On 17th February 2013, two rather modest flares were produced in quick succession. Both flares involve a highly sheared polarity inversion line (PIL). During the second flare (a GOES M1.9 class flare) an eruption was also produced. During the flare/eruption, the sunquake was detected in the positive magnetic field to the west of the highly sheared PIL. Figure 1 shows the location of the sunquake in the context of the coronal (SDO/AIA 171 Å and 304 Å wavebands) and magnetic field configuration (SDO/HMI line of sight data). The sunquake location is indicated on the HMI magnetogram by the red (6 mHz egression power) and blue (5 mHz egression power) contours.

The reconstruction of the coronal field

In order to probe the magnetic field configuration in the corona and investigate the likely sites of energy release and deposition, a non-linear force-free field (NLFFF) extrapolation was computed using a HMI vector magneotgram taken just over one hour before the flare/eruption using the method of Valori et al. (2010) [2]. The extrapolation revealed the presence of a flux rope with a hyperbolic flux tube (HFT) configuration above the PIL where the flare/eruption were produced and where the sunquake occurred. The magnetic field configuration could be decomposed into several key components. In increasing altitude, these components are listed below and shown in Figure 2:

  1. Magnetic field lines above the HFT (shown by orange field lines). These field lines correspond to an observed filament that erupted during the M-class flare. The filament location is shown in the upper left panel of Figure 2.
  2. A right-handed moderately twisted flux rope (shown by blue field lines) that is almost perpendicular to the polarity inversion line.
  3. Field lines that form an envelope surrounding the flux rope, which are essentially perpendicular to its central section (shown by green field lines).

Conclusions

The NLFFF model combined with the SDO observations leads us to conclude that the reconnection site responsible for the energy release that is ultimately responsible for the sunquake is located all around the top of the flux rope. This reconnection is triggered by the expansion of the erupting flux rope and its associated filament, liberating energy along the green field lines shown in Figure 2. Our investigation supports the idea that in this event the sunquake’s occurrence and location are strongly influenced by the magnetic field configuration. The magnetic field appears to act as a “lens”, focusing the released energy on to a particular site in the lower atmosphere. We have modelled with unprecedented detail the local magnetic environment in which the sunquake occurs, which can be then further used to test the available physical mechanisms that generated the energy transfer powering the sunquake. The magnetic field configuration is relevant to both wave and particle sunquake mechanisms, as the magnetic field acts as a guide for both. Such a particular magnetic lens configuration, as found in active region 11675, may not always be present in solar flare magnetic fields, which could explain the erratic nature of sunquake occurrence.

References

  • [1] Green, L. M., Valori, G., Zuccarello, F. P., Zharkov, S., Matthews, S. A., Guglielmino, S. L., The 2013 February 17 Sunquake in the Context of the Active Region’s Magnetic Field Configuration, 2017ApJ. 849, 40.
  • [2] G. Valori, B. Kliem, T. Török, and V. S. Titov. Testing magnetofrictional extrapolation with the Titov-Démoulin model of solar active regions. A&A, 519, 2010.

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84. The first NuSTAR microflare

Author: Paul Wright and Iain Hannah at the University of Glasgow.

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Introduction

Solar flares are observed to occur over many orders of magnitude, with smaller flares occurring considerably more often than large flares. This relationship is observed to be a power-law over many orders of magnitude [1]; however, it is not clear how small flare-like events can be and whether they share the same physical properties of large flares, both heating material and accelerating particles.

With the launch of the Nuclear Spectroscopic Telescope Array (NuSTAR) [2], Hard X-Ray (HXR; 2.5 − 78 keV) observations of faint, previously undetectable solar sources can be obtained. NuSTAR consists of two direct focusing telescopes and has over 10x the effective area, and a much smaller background than RHESSI. As NuSTAR is an astrophysics mission solar observations lead to various challenges [3], but recent studies have demonstrated that NuSTAR is highly capable of observing weak X-ray sources from the Sun. These observations range from occulted post-flare loops [4], to microflaring, and non-flaring active regions (ARs) [5,6,7], and quiet sun observations [8].

In this nugget we present the first joint observations of a microflaring active region with Hinode/XRT, SDO/AIA and NuSTAR. This equivalent A0.1 GOES class microflare (determined during the impulsive phase) within AR12333 shows heating of material to several million Kelvin, observed in Soft X-rays (SXRs) with Hinode/XRT, and was faintly visible in Extreme-UltraViolet (EUV) with SDO/AIA.

NuSTAR Imaging Spectroscopy

On 29th April 2015 Hinode/XRT was near-continuously observing AR12333, in which a small A0.1 microflare occurred, producing hot loops visible in the Hinode/XRT channels, as well as the hottest SDO/AIA channel (particularly the hotter Fe XVIII component of 94Å), see Fig. 1. During this time NuSTAR was producing two full-disk mosaics (its 12’x12′ field of view requiring at least 16 separate pointings to capture the whole Sun) and subsequently observed AR12333 four times, catching the impulsive phase of the microflare (see Fig 1.). The NuSTAR image of the microflare (inset Fig 2.) shows similar loop emission that matches that shown in Hinode/XRT and SDO/AIA. Crucially NuSTAR is an imaging spectrometer and so X-ray spectra are obtained, and fitted, for this region for each of NuSTAR’s four observations. During the pre-flare, decay, and post-flare phases, the NuSTAR spectra are well fitted by a single thermal model of 3.2−3.5 MK, but the spectrum obtained during the impulsive phase shows emission inconsistent with a such a model. We instead performed a two-thermal fit to this spectrum, finding a hotter component of about 10MK, which fits the bump in emission between 6 and 7 keV due to line emission from the Fe K-Shell transition [9]. This thermal fit, and no counts observed > 7 keV, indicate that NuSTAR is only detecting thermal emission from this microflare.

Microflare Multi-thermal Emission

With this microflare we can calculate the differential emission measure (DEM) distribution using SDO/AIA, Hinode/XRT, and NuSTAR, giving unprecedented coverage in temperature. By comparing the observed fluxes to those derived from the NuSTAR spectral fits, we find that the Hinode/XRT responses are about a factor of two too small, whereas the SDO/AIA Fe XVIII is as expected (see [5] for a complete discussion). Using these modified response functions we find the pre-flare DEM peaks at ~3 MK and falls off sharply by 5 MK; during the flare’s impulsive phase the emission above 3 MK is brighter and extends to 10 MK (Fig. 3). From these DEMs we determined the instantaneous heating rate during the impulsive phase of the microflare to be ~2.5 × 1025 erg s-1, consistent with the value obtained from the NuSTAR spectra [5].

Microflare Energetics

The NuSTAR spectrum from the impulsive phase of this microflare is purely thermal (Fig. 2). The non-thermal emission from accelerated electrons is either not present or too weak for NuSTAR to observe in this microflare. With this observation NuSTAR did have a limited spectral dynamic range due to a short effective exposure (about 2s from a 116s pointing) from a low detector livetime, one of the challenges of solar observations [3]. We can therefore determine upper-limits on the possible non-thermal emission. The upper-limits were calculated using the thick-target model with a single negative power-law, with index δ, above a low-energy energy cut-off, Ec (see [10]). Each upper-limit was obtained by iteratively reducing the model electron flux, until the X-ray emission was consistent with a null detection with NuSTAR. Fig. 4 shows these upper-limits for a range of power-law indices along with the three estimates for the thermal power (calculated as the instantaneous thermal energy during the impulsive phase minus that of the pre-flare time) plotted with 1σ uncertainties. The grey shaded region represents the required heating power, consistent with all three estimates. We find that for the accelerated electrons to be the source of the heating requires a power-law spectrum of δ = 7 with Ec < 7 keV, and a steeper spectrum would have a larger Ec.

Conclusions

This A0.1 microflare observed with NuSTAR shows heating of material up to 10MK, as well as non-thermal upper limits consistent with accelerated electrons being the source of this heating, though requiring a steep power-law spectrum down to low energies. This first NuSTAR microflare therefore strongly resembles much more powerful flares in terms of energetics. Another NuSTAR observation [6], has also shown a slightly smaller microflare behaves like a scaled down larger event. Future observations with NuSTAR during quieter, more optimal times should help probe other microflares, potentially detecting their non-thermal emission. The NuSTAR observations have shown the huge potential for highly sensitive solar X-ray observations and hopefully this will lead to more observations from missions like FOXSI, both in sounding rocket and possible satellite forms.

References

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83. Beam electrons as sources of Hα ribbons in a C-class flare

Author: Valentina Zharkova, Malcolm Druett and Eamon Scullion at the University of Northumbria.

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Introduction

Observations of solar flare onsets show a rapid increase of hard and soft X-rays, ultra-violet emission with large Doppler blue-shifts associated with plasma upflows, and Hα emission with red-shifts up to 1–4 Å [1, 2, 3]. Modern radiative hydrodynamic models account well for blue-shifted emission, but struggle to reproduce closely the red-shifted Hα lines observed at the flare onset. Here we present a joint hydrodynamic and radiative model showing that during the first seconds of beam injection the effects caused by beam electrons can reproduce Hα line profiles with large red-shifts closely matching those observed in a C1.5 flare by the Swedish Solar Telescope. The model also accounts for the timing and magnitude of upward plasma motion to the corona observed 29 s after the event onset in 171 Å by the Atmospheric Imaging Assembly/Solar Dynamics Observatory.

Observations of blue and red shifts during the flare onset

The flare occurred on 30th June 2013 in active region (AR) 11778 from 09:11 to 09:27 UT (Fig. 1a). AIA observations in 94, 171 and 304 Å channels presented in Fig. 1b show rather variable signatures. A bright, transient jet-like protrusion of plasma from the ribbon in the 171 Å AIA channel was detected between 15 and 29 s after the event onset, above a strong downflow regions in Hα emission (red contours) (Fig. 1b). The jet velocity in 171 Å, measured 29 s after beam injection at the location began, was 93 km s−1. At the same time, there are no jets seen in the 94 or 304 Å emission.

The Hα line observation sequence from 09:15.54 UT to 10:17:18 UT was made at the SST using the CRisp Imaging Spectro-Polarimeter (CRISP) [4]. The resulting Hα line profiles are shown in Fig. 2d, e. The CRISP observation captured the onset of a strong chromospheric downflow in the second ribbon area highlighted by the blue box in Fig. 2a, which is shown enlarged for the Hα red wing (Fig. 2b) and line centre (Fig. 2c). The red wing enhancement started in the 09:15:54 UT frame (Fig. 2d, red line), increased between 09:15:54 and 09:16:01 UT and peaked at 09:16:01 UT (Fig. 2d, purple line), 7 s after the flare onset.

Hydrodynamic and radiative models

We applied the HYDRO2GEN codes, which differ significantly from the other standard radiative hydrodynamic codes used in the solar physics community such as FLARIX [5], RADYN [6] and RH [7], as outlined below:
1) In the hydrodynamic models we have employed [8] the initial conditions are those of the quiet Sun chromosphere in hydrostatic equilibrium with an isothermal temperature distribution (T=6700K) without a heated chromosphere, as opposed to the RADYN models that have an (ad hoc) heated quiet corona overlying either quiet (VAL-C) or flaring (VAL-F) chromosphere.
2) The beam heating function, and beam electron number densities are derived using the continuity equation approach [9, 10], which provides smooth heating functions from the top of the flux tube to the photosphere. This is in contrast to the flux conservation approach [11] which can lead to excessively high, sharply-peaked heating rates [12].
3) The non-thermal excitation and ionisation rates due to collisions between beam electrons and ambient hydrogen atoms are taken from the exact analytical solutions provided by Zharkova and Kobylinskii [13].
4) The hydrogen ionisation degree is governed by non-thermal ionisation by beam electrons from all levels of hydrogen atoms and maintained by radiative transfer in the Lyman continuum.

Figure 3 shows plots of electron kinetic temperatures (a), macrovelocities (b) and plasma number densities (c) as functions of column depth calculated as a hydrodynamic response of the ambient plasma to injection of a power-law beam with the initial flux of 1010 erg cm−2 s−1 and spectral index γ=3. The initial QS chromosphere density is indicated by the straight line in Fig. 3c. The flaring transition region is swept by the beam towards 3 × 1018 cm−2, with the flaring chromosphere extending to 8 × 1019 cm−2 followed by a flaring photosphere (Fig. 3). The upward motion of flaring plasma is reflected in the macrovelocity plot (Fig. 3b) showing evaporation of chromospheric plasma upwards with velocities of about 93 km s−1s to the newly formed corona at the column depths between 1017 and 1019 cm−2) (Fig. 3b, box 1).
At the same time, the beam energy deposition leads to formation of a low-temperature condensation in the flaring chromosphere seconds after beam injection begins (Fig. 3a, b box 2) with a slightly increased temperature up to 104 K. This condensation moves as a shock towards the photosphere and interior with velocities up to 55 km s−1 (Fig. 3b). The density of this shock is about 1013 cm−3 (Fig. 3c).

Interpretation of observations

The temperature profile evolution for the model between 5 and 100 s shown in Fig. 3a (see box 1) reveals that the plasma can be detectable in the temperature range of log ⁡T=5.2 to log ⁡T=6.05 at the depths of the low flaring corona. The 171 Å channel is the most sensitive to this range, compared to other available AIA channels. Moreover, the velocity range derived from the AIA 171 Å channel, averaged at 93 km s−1, closely resembles the predictions of the model of a hydrodynamic response to plasma heating by an electron beam for a given temperature range, as shown in Fig. 3a.

The hydrogen radiative simulations clearly show that in the first seconds after the beam onset Hα line profiles are dominated by non-thermal ionisation by the beam electrons and the downward motion of the shock (see Fig. 3b, box 2). For this flaring event the beam has a relatively low initial energy flux about 1010 erg cm−2 s−1 resulting in a moderate increase of the Hα wing intensity (Fig.… continue to the full article

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