Author: Richard Morton Northumbria University.

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Introduction

Magnetohydrodynamic (MHD) waves are thought to play some role in the heating of cool, solar-like stars’ atmospheres, and acceleration of their winds. The role of acoustic (or longitudinal MHD) modes in energising stellar coronae has long been discounted, both from theoretical and observational considerations. The acoustic modes are predominantly excited by the turbulent convection (and are known as p-modes). Moreover, they clearly play some role in the structuring and dynamics of the chromosphere [1], with their energy being deposited there as they shock due to the increasing temperature. However, relatively recent theoretical results have shown that, in principle, the p-modes may generate Poynting flux in the corona. Given favourable alignments between the magnetic field and gravity, the longitudinal mode converts first to a fast MHD mode at the equipartition layer, where gas pressure and magnetic pressure are equal, followed by a secondary conversion to an Alfvénic mode at the transition region [2,3];  with only a small fraction of the available p-mode energy needing to be converted to satisfy the coronal energy budget.

If this process was in action in the Sun’s atmosphere, we should expect the coronal Alfvénic waves to display some signature of this. Previous studies [4,5] have indicated the presence of enhanced power (or ‘bump’) in the velocity power spectra of Alfvénic fluctuations derived from the Coronal Multi-Channel Polarimeter (CoMP [6]). However, these were limited to a few coronal regions on a couple of dates. Given the global nature of the p-modes and the expectation that they leak into the atmosphere through magneto-acoustic portals [7] (largely located at the network boundaries), one would anticipate that the enhanced power should be present globally, and should be more or less insensitive to the large-scale changes in the solar magnetic field activity cycle. We thought it was time to have a look! (For full details see [8].)

The coronal ‘bump’

CoMP has collected a large amount of data over a 7 year period, performing spectroscopic measurements of the Iron XIII 10747 Å coronal emission line on a near daily basis (see a previous nugget). We decided that in order to provide a more convincing case for the role of p-modes, we needed to examine the power spectra both globally and also across the 11-year cycle. This required a systematic analysis of Doppler velocity power spectra taken from many different sections of the corona. The accessible CoMP archive currently only has data from 2010, on the rise to solar maximum in 2013/2014, until 2018. We chose 5 different days between 2012 and 2015 to cover the rise and fall of the sunspot cycle and also had access to a special data set taken back in 2005, close to the end of the previous activity cycle. While this does not provide even sampling over the entire 11 year period, the Sun undoubtedly experienced large-scale changes in the global magnetic field over the time-span of the data.

We spilt the corona into 5-degree bins in order to average the velocity power spectra and increase signal-to-noise. This is a coarse approach to the problem, neglecting the presence of different magnetic geometries in each bin (see for example the structures in Figure 1); different magnetic regions, e.g. active region, coronal hole, having power spectra with different properties [9]. Each average spectrum was then fit with two models, a power law and a power law with a localised enhancement of power (represented by a log-normal function), where the enhanced power represents the contribution of p-modes to the coronal Alfvénic wave power spectra. We compared the two models’ ability to describe the spectra through Information Criteria. The results found that in the large majority of cases (>95%), the power law with the enhanced power was found to be the better model. This was the same for each data set we examined. The signature of p-modes is everywhere in the corona, and always seems to be present

Moreover, the model parameters that described the enhanced power, namely the frequency associated with the location parameter and the frequencies associated with the scale parameter are remarkably similar across the corona and all dates (Figure 2). The location of the enhanced power is predominantly around 4 mHz, shifted towards higher frequencies than the characteristic p-mode frequency of 3.3 mHz, while the typical scale was just greater than 1 mHz.

The prominent nature of this enhanced velocity power signal in the Sun’s corona is apparent, as it is still visible even after averaging over all the coronal power spectra (Figure 3). This signal is clearly global and ever present, and our results indicate that it identifies a fundamental component of the corona – namely the p-mode ‘driven’ Alfvénic waves (the quotes around driven are meant to imply that it is not directly the p-modes that drive the waves).

Conclusions

What does it mean to have p-mode ‘driven’ Alfvenic waves in the corona? One major implication is that many models of coronal heating and wind acceleration by Alfvénic wave turbulence are ignoring a sizeable source of Poynting flux. In the current models, it is largely assumed that the horizontal buffeting of photospheric flux tubes is the sole source of Poynting flux, with the waves entering the corona from below having relatively high frequencies (< 4mHz). Hence, we are suggesting that these models are missing a key piece of physics! Moreover, in order to understand the heating/acceleration of solar plasma requires including the extra energy injected at the transition region via the mode conversion process. While the broader consequences of including this extra energy are unknown, it will certainly play a role in energising the corona.

On a more speculative note, our results may provide a way to observe Alfvénic waves on other cool, magnetised stars, namely ones that generate stellar oscillations (of which Kepler has found many [10]). Current instrumentation is still some way from providing detailed observations of other stars’ coronae, although velocity measurements have been made in Antares’ atmosphere [11].… continue to the full article

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Author: David Kuridze at the Aberystwyth University, Mihalis Mathioudakis at the Queen’s University Belfast, Huw Morgan at the Aberystwyth University.

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Introduction

The structure and dynamics of the most energetic events in the Sun’s outer atmosphere such as flares, eruptions, coronal loops and coronal mass ejections are controlled by the magnetic field. Coronal loops are the fundamental building blocks of the coronal field – these are bright, dense structures that trace field lines that connect photospheric regions. Measurements of the magnetic field in loops is key to our understanding of the corona and is crucial to solve the long-standing problem of coronal heating. Despite developments in observation and analysis, a direct quantitative measurement of the magnetic flux density of coronal loops remains a central problem in solar (and stellar) physics [1].

Flare loops

Despite its promise, solar coronal polarimetry is extremely difficult due to the inherently low signal-to-noise ratios [2,3]. Current instrumentation can only achieve high-resolution polarimetric measurements during certain favorable conditions. In September 2017, Active Region (AR) 12673 produced the most powerful flares of solar cycle 24 as it was rotating from disk centre to the limb. On September 10 2017, the AR was just behind the west limb when it produced an X8.2-class flare at 16:06 UT. High spatial and temporal resolution observations of the flare were acquired with two instruments – CHROMospheric Imaging Spectrometer (CHROMIS) and CRisp Imaging SPectropolarimeter (CRISP) on the Swedish Solar Telescope (SST) in La Palma, Spain. The dataset consists of spectral imaging in Hβ (4861 Å) and full Stokes spectropolarimetry in the magnetically sensitive infrared Ca II line at 8542 Å (Landé g factor 1.1). These lines are usually formed under chromospheric conditions where the plasma temperature is ~7 000 – 20 000 K. The X8.2 flare led to intense heating and evaporation of dense chromospheric plasma into the loop system. The chromospheric plasma underwent rapid cooling (condensation) and the coronal loops subsequently filled with cool plasma radiating strongly in chromospheric lines including Ca II 8542 Å (Fig.1 & 2). The off-limb location of the AR minimized any contamination caused by line-of-sight effects, providing polarisation data of unprecedented quality.

Magnetic field map

Circular polarisation images (Stokes V) of the observed flaring loops acquired in Ca II 8542 Å reveal strong polarization signals along the loops (Fig. 2). The Stokes V component of the polarized light is related to the first derivative of the intensity profile (Stokes I) and the line-of-sight (LOS) component of magnetic field (B_LOS) through the weak-field approximation (WFA) [4]. The WFA is applicable when the Zeeman splitting of a line is much smaller than its Doppler broadening, and the magnetic field and LOS velocity are uniform along the LOS. These requirements are satisfied in the observed flare coronal loops [4]. We use the WFA to produce maps of the LOS magnetic field component (B_LOS) for the flare coronal loops (Fig. 2). The histograms in Fig. 2 are the distributions of the magnetic field for 3 different height ranges. B_LOS of the loop apex region (layer a, between 18 and 26 Mm above the solar surface) ranges from 50 to 180 G with median 90 G. The corresponding values at mid-heights (layer b, 9 to 18 Mm) is as high as 300 G with median 140 G. The uncertainty of the B_LOS in Fig. 2 is less than 20% [4].

The spectra are fitted to Gaussian functions to determine the bulk plasma LOS velocity (Fig. 3). Mapping this value (Dopplergrams) reveals that the left and right part of the loop structures contain regions of red- and blue-shifted plasma, respectively, with velocities between ~ 10-35 km/s (Fig 3A). Movies of this event and time-distance diagrams show strong downflows of dense plasma from the loop apex toward the footpoints. These velocities are the two orthogonal components of the downflowing plasma motions with respect to the observer – the LOS component (Dopplergrams) and the perpendicular component (apparent) . Since the moving plasma is confined to the loops, the direction of the flow velocity follows the direction of the magnetic loops. The average viewing angle of the magnetic field and velocity calculated from the ratio between the LOS and perpendicular components of the velocity vectors is approximately ~ 60 – 80 degrees (Fig. 3). This yields a median of total magnetic field strength, of around 350 G at the loop system apex (region a in Fig. 2), and 420 G at mid-heights (region b).

Conclusions

This unique observation of flaring coronal loops at the solar limb, using high-resolution imaging spectropolarimetry, has allowed us to determine their magnetic field with unprecedented accuracy due to the vantage point, orientation and nature of the chromospheric material that filled the flare loops. We find coronal magnetic field strengths as high as 350 Gauss at heights up to 25 Mm above the solar limb. These measurements are substantially higher than a number of previous estimates and could have considerable implications for our current understanding of the extended solar atmosphere. This constraint is crucial for physical models of coronal ARs, flares and eruptions, and can be used to provide a validation of widely-used numerical methods for the extrapolation of photospheric magnetic fields in the corona. Furthermore, the result is important for upcoming new-generation ground-based solar telescopes such as the 4 m Daniel K. Inouye Solar Telescope (DKIST) and the European Solar Telescope (EST). These telescopes will have advanced chromospheric polarimetric capabilities, which, as demonstrated here, can provide powerful diagnostics for the coronal magnetic field.

References

• [1] Casini, R., White, S. M., & Judge, P. G. 2017, SSRv, 210, 145
• [2] Lin, H., Kuhn, J. R., & Coulter, R. 2004, ApJL, 613, L177
• [3] Lin, H., Penn, M. J., & Tomczyk, S. 2000, ApJL, 541, L83
• [4] Kuridze, D., Mathioudakis, M., Morgan, H., et al. 2019, ApJ, The Astrophysical Journal, 874, 126

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Authors: Iain Hannah and Lyndsay Fletcher
University of Glasgow.

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Introduction

Welcome to the 100^th UKSP nugget*. To mark this milestone we thought it was time to have a look back at some of the solar physics research we have been highlighting through the nuggets since the first one back in July 2010. Thanks to all the authors for their work and a visual summary of all their nuggets is shown in Figure 1.

What have the nuggets been about?

What topics do the UKSP Nuggets cover? A word cloud of all the nugget titles is shown in Figure 2. It appears that “coronal” and “corona” are our favourite part of the Sun, “magnetic”, “reconnection” and “waves” are also well covered, but not so many people are working on “sunspots” or “chromosphere” or “particles”, it seems! If you think that your favourite solar word is too small here, why not write a nugget?

UKSP nuggets are written by authors throughout the UK, as can be seen in Figure 3. Most of the main centres for solar physics in the UK are covered – can you name them? If you have a piece of solar physics work that is led from the UK and want to write a nugget for us, please get in touch. General advice is still under 1,000 words is best if you want the most readers, however from the previous nugget summary the relationship between page views and length is no longer as clear, see Figure 4.

Who is reading the nuggets?

Each nugget still typically gets about 100 views just after publication and that slowly grows over time, with a long term average of unique views per nugget at nearly 400. Some of the most popular nuggets have several thousand views. In total we have had nearly 40,000 page views, and 34,000 unique pages views of the nuggets. The readers of the nuggets are predominantly from the UK (43%), but we have page views from around the world, see Figure 5. The rest of the top ten countries for views are United States (15%), China (5%), India (5%), Russia (4%), Mexico (3%), Poland (2%), Germany (2%), Ireland (2%).

What next?

UKSP nugget 101 of course. Solar physics continues to be a highly productive field in the UK, which will no doubt continue given our leading involvement in new telescopes and missions like Solar Orbiter, DKIST (and EST), SKA etc, as well as the latest numerical and analytical work. We are very grateful to everyone who has contributed to the UKSP nuggets, and are always looking for new UK-based nuggeteers (particularly early-career researchers) to help tell the world about our solar physics research. We would also like to hear your comments and suggestions about the UKSP nuggets.

*Technically this is UKSP Nugget #101, as a previous summary at the 3 year mark was a special unnumbered nugget.… continue to the full article

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Authors: Natasha Jeffrey, Lyndsay Fletcher^* and Nicolas Labrosse at the University of Glasgow and Paulo Simões at MacKenzie Presbyterian University, São Paulo.

^*Also at Rosseland Centre for Solar Physics, University of Oslo

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Introduction

From the optically thin lines emitted by hot plasmas it is possible to obtain properties such as density, temperature and flow speed. If a spectral line is formed in thermodynamic equilibrium in a plasma with a Maxwell-Boltzmann distribution of speeds then the line width can also be used to deduce the temperature of the emitting ions. However, in some cases, the line width is observed to be significantly larger than expected for a plasma with temperature consistent with the ions present [1]. This broadening is generally called “non-thermal broadening”. It could have different causes, but in this nugget we present observations made before and during a solar flare, in which the non-thermal broadening is consistent with turbulent motions in the emitting plasma, which decrease as the flare intensity increases. This is evidence for transfer of energy from plasma turbulence to heating – i.e. the kinetic energy of particles in the transition region. Evidence of turbulent energy transfer has already been seen during a flare in a coronal source [2] but here we present observations of broadening in the lower atmosphere, which clearly precede the flare brightening.

IRIS observations of SOL2016-12-06T10:36:58

The IRIS spacecraft observed a B-class flare at 1.7s high cadence, in sit-and-stare mode. This is the highest cadence yet for such observations. Figure 1 shows the location of the IRIS spectrometer slit on the background of the IRIS 1400 Å slit-jaw image. The white contours indicate the SDO/AIA 131 Å footpoints and the pink contours are the RHESSI 6-12 keV source. No higher energy RHESSI emission could be imaged.

The Si IV line at 1402.77Å was fitted with a Gaussian profile to obtain the line intensity, line-of-sight speed and line broadening. These values are interpreted under the assumption that Si IV, emitted at temperatures close to 80,000 K, is optically thin. Modelling suggests that this might not not be the case in some flares [3] but this event was weak enough that the optically thin assumption holds.

Time-dependent line broadening and flows

The results of the spectral fitting are shown in Fig. 2 below. In the upper and lower panels, the grey shaded area shows the intensity in Si IV, integrated over 2″ along the slit. About 2 minutes of data are analysed, and the strong flare brightening lasts less than a minute at this wavelength. Superposed in the top panel is the FWHM, or non-thermal velocity (RH axis), as well as the RHESSI 6-12 keV flux. On a timescale of around 10 seconds, the line broadening increases strongly, and does so well before the flare brightening starts. As the Si IV intensity increases, signifying atmospheric heating, the non-thermal width decreases. The non-thermal width then has another 2 peaks, separated by about 11s, before decaying back to pre-flare values.

If we interpret the non-thermal broadening as due to turbulence in the transition region then the implication of this observation is significant. The timing suggests that turbulence in the region starts well before the flare heating, rather than occurring as a result of heating. Instead, the heating only happens after the turbulent broadening peaks and as it decays. This is a strong indication of cause and effect – energy is transferred from the turbulence to the particles in the transition region plasma.

The lower panel of Figure 2, plotting the Si IV line centroid, shows that the peak in the turbulence occurs at the same time as a strong blueshift. Normally, blueshifts would be associated with chromospheric evaporation caused by flare heating, but here the blueshift occurs before the flare heating indicated by Si IV intensity. So we can rule out chromospheric evaporation as a cause. However, we have found in simple modelling that the patterns of line broadening and blueshift together can be explained by traveling waves.

Using a simple model of motion of fluid roughly perpendicular to the field direction, caused by traveling Alfvénic waves, we have shown that both the line broadening and the line shift can be approximately reproduced by a superposition of a small number (around 10) of traveling waves, with an amplitude and phase varying with time, and a wavelength comparable to the thickness of the Si IV emitting region. Figure [3] shows the plasma velocity as a function of position, due to the modelled waves, with colour representing time (top panel). Then the resulting velocity amplitude as a function of time at a randomly chosen cut is shown in the middle panel, and the line broadening and line shifts that this would generate at the bottom.

The basic properties of the velocity observations are reproduced, i.e. an increase in the line broadening followed by a small number of oscillations, with the line centroid simultaneously blueshifted. A spectrum of waves entering the transition region would be very likely to reflect, interact and lead to a turbulent cascade ending in heating. How these waves are generated is, however, another question.

Conclusions

Very high cadence IRIS observations have allowed us to detect and characterise the evolution of non-thermal line broadening, interpreted as a signature of turbulence, in the Si IV transition region line before and during a small flare. The relative timing of the onset of line broadening and of flare heating indicated by the rise in the Si IV intensity shows clearly that the heating does not cause the turbulence; if anything it is the other way around. A possible interpretation is that the flare launches magnetic disturbances from the corona towards the solar chromosphere, carrying energy that ultimately dissipates in the form of ion heating as the waves travel through the transition region, setting up a turbulent cascade. This may provide a significant mode of energy transport in flares. The full results of this study are described in [4].

References

• [1] E. Antonucci, M.

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Authors: Alexander James, Lucie Green, Gherardo Valori, and Lidia van Driel-Gesztelyi at MSSL UCL.

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Introduction

To identify the mechanisms that initiate Coronal Mass Ejections (CMEs), we must first know what the magnetic structure of the solar corona is prior to eruptions. Measurements taken at 1 AU reveal that at least some — or perhaps all — interplanetary CMEs contain magnetic flux ropes: bundles of twisted magnetic field lines [1,2]. However, there is still debate regarding the pre-eruptive configuration of CMEs: do these flux ropes form before the onset of eruption [3], or after the onset of eruption of a sheared magnetic arcade [4,5]? In this work, we use a combination of EUV observations and magnetic field reconstruction to investigate the pre-eruptive configuration of a solar active region and to identify the mechanism that caused its eruption.

A Flux Rope in the Solar Corona

An interplanetary CME with a flux rope structure was identified in measurements from the Wind spacecraft at 1 AU and found to have originated from NOAA active region 11504 on 14 June 2012 [6]. At that time, the source active region was close to the centre of the Sun as viewed from the Solar Dynamics Observatory (SDO), meaning reliable measurements of the surface magnetic field in the active region were available.

During the two hours leading up to the eruption, an S-shaped plasma emission structure called a sigmoid was observed in the 131 Å extreme-ultraviolet (EUV) channel of SDO/AIA (Figure 1a). A sigmoid is a telltale signature of a magnetic field with at least one turn of twist along its length [7], suggesting that a flux rope was present in the active region during this time. The sigmoid appeared above a bright flaring arcade, indicating that the flux rope was forming high up in the corona via magnetic reconnection, a process that can change the local configuration of the magnetic field via the extremely small but finite resistivity of the coronal plasma. This is not the most common location where flux ropes are formed, but a number of other observations helped build the case that a flux rope may have formed high in the corona of the active region at least two hours before the eruption, including twin EUV footpoint dimmings, flare ribbons, radio emission, and spectroscopic composition data from Hinode/EIS [8].

However, no single observation provides definitive proof that a flux rope was present. As a next step, a model of the active region magnetic field was produced by extrapolating the surface magnetic field measurements under the assumption that the coronal field is free from Lorentz forces (force-free). The extrapolated coronal field indeed contained a flux rope that extended high up in the corona and matched the observed sigmoid and footpoint signatures remarkably well (Figure 1b) [9].

With the validity of the extrapolated coronal magnetic field confirmed against observations, we could study it quantitatively. The decay index quantifies how rapidly the overlying stabilising field strength changes with height. In idealized theoretical models, a semi-toroidal flux rope would be unstable to expansion under the hoop force if the decay index is larger than n_crit=1.5 at its axis [10,11]. The extrapolation reveals that the flux rope lies in a region where the decay index is comparable to this threshold, meaning that the eruption of the flux rope is likely to have been driven by this so-called torus instability.

As for how such a coronal flux rope was formed, we observed that several hours before the eruption, the active region contained two sets of J-shaped magnetic loops that underwent magnetic reconnection in the corona and transitioned to the flux rope and underlying flaring arcade. Magnetic flux emergence was ongoing in the active region during this time, and subsequently-emerging fragments of magnetic flux were seen to ‘orbit’ around pre-existing flux (see the movie in Figure 2). We argue that the footpoints of the observed J-shaped loops were rooted in these sunspot fragments that wrapped around each other, providing a triggering mechanism for the sheared coronal field to converge, reconnect, and build the flux rope in the corona.

Conclusions

In this case study, we observationally inferred the presence of a magnetic flux rope that formed in the solar corona at least 2 hours before its eruption. This scenario was confirmed with an observationally-constrained model of the active region, produced using the nonlinear force-free field extrapolation of surface magnetic field measurements. According to our analysis of the extrapolated coronal model, the eruption was driven by the torus instability. Finally, the observations suggest that the motion of newly-emerged magnetic flux in the photosphere enabled reconnection in the corona that built the flux rope.

References

• [1] Gosling 1990
• [2] Marubashi et al. 2015
• [3] Green & Kliem 2009
• [4] Moore et al. 2001
• [5] Antiochos et al. 1998
• [6] Palmerio et al. 2017
• [7] Green et al. 2007
• [8] James et al. 2017
• [9] James et al 2018
• [10] Bateman 1978 Cambridge, Mass., MIT Press. 270 p.
• [11] Kliem & Török 2006

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Author: Viktor Fedun, Gary Verth at the The University of Sheffield, J. J. González-Avilés at the IGUM-UNAM, F. S. Guzmán at the IFM-UMSNH, S. Shelyag at the Deakin University and S. Regnier at the Northumbria University.

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Introduction

There is evidence that Type II spicules are due to magnetic reconnection [1-4], oscillatory reconnection [5,6] or more recently it was shown in [7] that spicules occur when magnetic tension is amplified and transported upward. In this work, using 3D numerical simulation of the reconnection process in the photosphere-corona region, we show the formation of a jet with characteristics of Type II spicules.

In our model we assume (i) a completely ionised solar atmosphere which is governed by the resistive MHD equations subject to a constant gravitational field, and (ii) the solar atmosphere based on the C7 model in combination with a 3D potential magnetic field, extrapolated up to 10 Mm above the photosphere, with photospheric values taken from a large-scale, high-resolution, self-consistent simulation of solar magneto-convection with MURaM [8,9]. With these ingredients at the initial time we evolve the system using the resistive MHD equations on the numerical domain x∈[0,6], y∈[0,6], z∈[0,10] Mm, covered with 240×240×400 grid cells.

Jet formation and evolution

The 3D evolution of the magnetic field and the temperature on the plane x=0.1 Mm are shown in Figure 1. The corresponding 2D slice at the plane x=0.1 Mm is plotted in Figure 2. Notice that at time t=15 s the jet starts to develop at the transition region level z≈2.1 Mm where there is a strong current density, which is an indication of reconnection happening. By t=60 s the jet has features of a Type II spicule, with a base located at z≈2 Mm and a height of about z≈7 Mm measured from the transition region (see Figure 2). This in agreement with the observed heights of the Type II spicules, between 3-9 Mm. Finally at time t=90 s the spicule reaches the 9 Mm height.

In order to locate regions where magnetic reconnection has occurred, we show a 2D slice of the evolution of |J|(A m^-2) and temperature contours in Figure 3. At time t=60 s, when the spicule is well formed, a strong current is located around (y,z)∼(2,2) Mm, at the basis of the spicule, which is consistent with the results in Figure 2. At t=90 s, the regions of strong current are still located at the bottom of the spicule. This analysis shows that magnetic reconnection mainly happens at the chromosphere and the transition region.

We compare the forces due to the magnetic field and gas pressure. The results of the evolution of the ratio |J×B|/|∇p| and temperature contours (K) are shown in Figure 4 in the plane x=0.1 Mm of the 3D domain. Notice that at time t=15 s, which is the time when the spicule starts to form, the Lorentz force dominates. At time t=60 s, the Lorentz force is stronger exactly where the spicule forms. This dominance is still clear at t=90 s. This analysis shows that Lorentz force is an important ingredient of the jet evolution.

Conclusions

Based on a 3D numerical simulation of a small region in the solar atmosphere, which spans from the photosphere to the corona, we show the formation of a jet structure with characteristics of a Type II spicule, specifically the morphology, upward velocity range and time-scale of formation. This result provides a simple explanation and is in contrast with that in [7], where out of 2D simulations the formation of spicules is explained in terms of the amplification of the magnetic tension and the interaction between ions and neutrals. In our simulation we show that the full Lorentz force is important in the process of formation, which is consistent with the results obtained in the simulations of twisted magnetic flux tubes [10] and in the formation of solar chromospheric jets [11]. Our findings also show that vorticity motions are important for jet formation. By analyzing the velocity field in a specific cross-cut of the spicule we can track the circular displacement of plasma that eventually can be identified with blue-red shifts. We invite you to read [12] to review the analysis in more detail.

References

• [1] Isobe, H., Proctor, M. R. E., & Weiss, N. O. 2008, ApJ, 679, L57
• [2] De Pontieu, B., McIntosh, S., Carlsson, M. et al. 2007, Science, 318, 5856
• [3] Archontis, V., Tsinganos, K., & Gontikakis, C. 2010, A&A, 512, L2
• [4] González-Avilés, J. J., Guzmán, F. S. & Fedun, V. 2017, ApJ, 836, 24
• [5] Heggland, L., De Pontieu, B., & Hansteen, V. H. 2009, ApJ, 702, 1
• [6] McLaughlin, J. A. et al. 2012, ApJ, 749, 30
• [7] Martínez-Sykora, J. et al. 2017, Science, 356, 1269
• [8] Shelyag, S., Mathioudakis, M., & Keenan, F. P. 2012, ApJ, 753, L22
• [9] Vögler, A., Shelyag, S., Schüssler, M., et al. 2005, A&A, 429, 335
• [10] Kitiashvili, I. N. et al. 2013, ApJ, 770, 37
• [11] Iijima, H. & Yokoyama, T. 2017, ApJ, 848, 1
• [12] González-Avilés, J. J. et al. 2018, ApJ, 856, 176

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Author: Gareth Hawkes and Mitchell Berger at the University of Exeter.

Introduction

Magnetic helicity is a measure of how twisted and interlinked the magnetic field is [1], given as the integral over all (relevant) space of the product of the magnetic vector potential and magnetic field. As the poloidal field lines emanating from the photosphere are wrapped around each other by the differential rotation at and below the solar surface, helicity increases directly in line with the amount of toroidal field being captured for a following cycle’s sunspots. Here, poloidal field refers to that which has a component normal to the surface, whereas the toroidal field does not. We theorise that the magnetic helicity flux during a solar minimum and the strength of the following cycle (determined by the sunspot number) should be well-aligned.

In the video below, we plot a potential field source-surface (PFSS) model of the magnetic field out to two solar radii, over the total time period used in the following study. Between solar cycles, the magnetic field is predominantly polar (with its polarity being flipped during each cycle). This polar field slowly descends over a solar cycle, forming much more complicated structures.

Magnetic Helicity Flux predicting Solar Activity

The polar field strength of the magnetic field has been used as a strong predictor of the upcoming solar cycle [2], amongst other techniques. Differential rotation wraps the polar field around the axis of the sun, creating the toroidal field seen in sunspots. The most fundamental attempts employ flux transport to model a simplified version of the solar dynamo (e.g. [3]). However, slight changes in the model parameters can result in wildly different results, which can be seen by comparing to other similar models [4]; while one model [3] made an accurate prediction of the strength of Solar Cycle 24, the other [4] made an under-prediction. In these models, there is also a large dependence on initial conditions.

Magnetic helicity flux is given by the alignment of the magnetic field’s vector potential with that of the velocity field on the photosphere, multiplied by the magnitude of the radial magnetic field as a scale factor. Helicity injection results in a stronger toroidal field, with a measureable coefficient describing the degree to which the poloidal field has wrapped around the toroidal field. Taking our velocity field as the differential rotation profile only gives a measure of the helicity injection or ejection due to the effect of differential rotation, or in other words how wrapped our field lines have become. This is equivalent to increasing the strength into the following cycle’s toroidal field. Whilst polar field does make up the majority of helicity injection, it disregards helicity generation in the sub-polar regions. Thus, magnetic helicity should be an as strong, if not stronger, predictor of the upcoming solar cycle as compared to at least polar field.

In Figure 2 we give the magnetic helicity flux in the northern hemisphere against total sunspot number for the period 1976-2018. Magnetic helicity is calculated using magnetic field data provided by the Wilcox Observatory, in the form of spherical harmonic co-efficients of a Potential Field Source Surface model. The velocity field representing differential rotation is taken from an analytical formula. Normalisation is performed relative to the highest peak in an individual data set. There is a clear correlation over the first two cycle pairs (a cycle pair refers to a helicity cycle followed by its closest sunspot cycle), as predicted. The weakest correlation is in the most recent cycle, where we saw an abnormally weak level of solar activity at the beginning and during the Solar Cycle.

Does helicity flux out-perform polar field? In Table 1, we give the ratio of the integrated values of magnetic helicity flux and polar field with respect to sunspot number. Here, polar field is taken as the absolute value of the average radial field over a fifteen degree polar cap. Naturally, integrating helicity flux gives a value of the amount of helicity which has been injected/ejected in time t. Integrating over sunspot number is less physically meaningful, but should be representative of strength when we consider that sunspots which are counted over two consecutive data points are likely to be more intense due to their longevity. These ratios should then correspond to how well the “amounts” of helicity and polar field generated during a cycle correlate with the “amount” of sunspot activity. We see a clear preference for helicity flux, where the ratio is consistently higher.

If we properly entertain the idea that the dynamos governing the northern and southern hemispheres could be imperfectly coupled or out of sync, we achieve much better results for the most recent solar cycle, the results of which are shown in Figure 3. Here, we plot magnetic helicity flux versus sunspot number, both of which are now constrained to the their respective hemispheres: North in the top panel and South in the bottom. This data set is limited to 1990-2018, due to a limitation in data availability. Importantly, the weakest (most recent) cycle pair are now much more closely aligned in terms of their normalised amplitude, for the Northern hemisphere. In the South, we see a stronger alignment of the 1990 cycle.

Conclusions

We conclude that magnetic helicity appears to be a stronger predictor of solar activity than polar field strength. This may be due to the association of helicity flux with the generation of toroidal field which governs the next solar cycle. Additionally, we found evidence that supports the theory that the dynamos in the north and south are either imperfectly coupled or at minimum out of phase with each other. We invite you to read [5] for a more detailed analysis.

Acknowledgements

The Wilcox Solar Observatory data used in this study to calculate magnetic helicity flux were obtained via the web site wso.stanford.edu, courtesy of J.T. Hoeksema, with thanks. Sunspot Number Source: WDC-SILSO, Royal Observatory of Belgium, Brussels. G. Hawkes would like to thank the STFC for their funding under grant ST/N504063/1.… continue to the full article

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Author: Patrick Antolin, Ineke De Moortel at the University of St Andrews, Don Schmit, Bart De Pontieu at LMSAL and Tiago M. D. Pereira at the University of Oslo

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Introduction

Spicules are one of the most active research subjects in solar physics, but also one of the most thorny. These chromospheric structures correspond to cool and dense jets of plasma that protrude into the corona, often in a ballistic trajectory. They are thought to constantly refill coronal loops with plasma. Because they are everywhere across the solar surface, and because they are highly energetic, their importance has been highlighted for studies of both the chromosphere and the corona, making them also a subject of great debate.

Of particular relevance is the unsolved problem of how the corona is heated, and an attractive theory is that spicules funnel most of the energy into the corona and that the heating mechanism already operates in the spicule [2].

The high energies of spicules manifest in their large swaying motions with amplitudes of a few 10 km/s, which suggest an Alfvénic wave driver with large enough energy flux up into the corona (on the order of 10⁶ erg cm^-2 s^-1). Alfvénic waves are transverse MHD waves (whose main perturbation is transverse to the magnetic field) who have as main restoring force the magnetic tension force. Among these waves is the kink wave, which is characterised by being able to displace the entire flux tube in which the spicule exists [3].

The swaying motion of the kink wave induces shear flows at the edges of the flux tube that are expected to lead to Kelvin-Helmholtz instabilities [4,5]. The resulting vortices are known as Transverse Wave Induced Kelvin-Helmholtz rolls (TWIKH) and are very efficient at mixing the plasma external to the spicule with the internal plasma [6]. Furthermore, the TWIKH rolls generate a turbulent-like state with a cascade of the wave energy to smaller scales and eventual dissipation in the myriads of current sheets and vortices [7]. TWIKH rolls, therefore, constitute a potential dissipation mechanism for kink waves that could play a role in coronal heating.

Observations

In this work [1] we analysed coordinated observations of an off-limb region near the north pole of the Sun (see Figure 1) made with Hinode/SOT imaging in the Ca II H spectral line that forms at 10,000-20,000 K, and with IRIS, an imaging spectrograph, here observing in the Mg II k line formed at 15,000 K and Si IV line formed at 100,000 K. This coordinated observation allowed a high resolution and multi-temperature view of several spicules during their evolution.

The spicule shown in Figure 1 highlights one of the main characteristics of the spicule: it is multi-stranded when observed at high resolution (SOT view), while it appears monolithic in the IRIS channels. The ensemble of strands sways transversely and in-phase, suggesting a kink mode. The multi-stranded nature is particularly clear during the swaying.

Thanks to the location of the IRIS slit (white horizontal line in Figure 1) we could record the spectral signatures during the swaying. Note in Figure 2 that during the transverse oscillation large variations in intensity are observed in the chromospheric lines. The Doppler shifts change sign at times of maximum displacement. Also, the location across the spicule where the transition occurs appears ragged. Furthermore, the line widths increase during the oscillation.

Simulations

In order to test the role of the kink wave on the observed signatures, we performed 3D MHD simulations combined with radiative transfer and optically thin modelling of a spicule-like structure undergoing a kink oscillation. The spicule consists of a density-enhanced (x50 -> 6×10¹⁰ cm^-3), colder (1/100 -> 10⁴ K) and low-beta structure with respect to a coronal background. There is a smooth transition between the external and internal plasma, with a slightly enhanced internal magnetic field to keep hydrostatic balance. The spicule is connected upwards with a coronal loop through a transition region (see Figure 3).

The spicule oscillates in the fundamental kink mode with an amplitude of 12 km/s in the top part of the spicule, and a period of 255s. After half a period TWIKH rolls are produced that deform and locally twist the spicule. Strong currents are produced that propagate as Alfvénic perturbations. The TWIKH rolls produce density-enhanced regions that appear as strands in the synthesised chromospheric lines. Because of the collective nature of the kink mode, the strands oscillate mostly in-phase (see lower panels of Figure 2), but a ragged transition in the Doppler shifts is generated because the TWIKH rolls lie mostly at the surface of the spicule, where strong backward flows exist (fuelled by resonant absorption). The Doppler maps also appear stranded because of the distinctive velocities of the TWIKH rolls. Furthermore, since rolls exist at multiple scales (with different dynamics), high-frequency Doppler shifts and increased line widths are produced.

The Ca II H source function (see Figure 3) in this model has an annular shape. Because of the strong density and temperature variations at the edge of the spicule due to the TWIKH rolls, this source function varies strongly, generating a bursty intensity evolution for this line. On the other hand, the mixing from the TWIKH rolls and the wave dissipation is not enough to produce an increase in intensity in hotter lines as is usually observed.

Conclusions

Despite our very simple spicule model, we have shown that kink modes lead to TWIKH rolls that can successfully explain several aspects of spicules:
– Multi-stranded structure
– In-phase oscillation of strands
– Large chromospheric intensity variation (particularly in the Ca II H line)
– Ragged Doppler shift sign changes at maximum displacement
– Increased line widths

On the other hand, our model fails to increase the spicule temperature enough to account for the observed increase in transition region and coronal lines. This is likely due to the simplicity of our model and suggests other mechanisms at work [8].

References

• [1] Antolin, P.,

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Authors: Tim Horbury and David Stansby at Imperial College London and Lorenzo Matteini at LESIA, Paris.

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A re-analysis of old solar wind measurements, taken near the Sun over 40 years ago, has revealed the presence of short velocity “spikes” in the plasma, which can take its speed to over 1000 km/s. These discrete spikes carry a significant fraction of the total momentum and energy of the wind and have striking similarities to results from recent simulations of coronal jets, helping to explain the energy budget of the solar wind and raising the exciting possibility that we will directly measure these jet outflows with the upcoming Parker Solar Probe and Solar Orbiter missions.

Introduction

How does the Sun produce a wind, flowing at several hundred km/s, that fills interplanetary space? The solar wind is a hot, near-collisionless, nearly fully ionised plasma. We know the basics: that the hot corona, being at a far higher pressure than interstellar medium, accelerates away until it is super-sonic and super-Alfvenic, carrying with it the Sun’s magnetic field. The fastest, smoothest wind comes from coronal holes at around 750 km/s and it has generally been assumed that this is a steady-state process. Models (see [1] for a recent review) of the wind cannot achieve these speeds without the addition of large amplitude Alfven waves, which help to drive the plasma outward. Theories of the origins of the Alfven waves in fast wind include footpoint motion driven by granular overturning, and broadband waves in the corona – which, if either, is correct?

Solar wind velocity spikes

We are about to embark on a renewed exploration of the inner heliosphere, with the Parker Solar Probe and Solar Orbiter missions. In preparation for the new data, the group at Imperial has been re-examining, using new methods, the closest measurements we have to the Sun, from the twin US-German Helios spacecraft in the 1970’s, see Figure 1. We have fully re-processed the proton measurements, and recently released the first publicly available reliable data set [2].

Our work on high speed solar wind streams at just 60 solar radii [3] has revealed a population of short, large amplitude velocity increases (Figure 2). While we are limited by the 40s resolution of the plasma data, it is clear that these events are common (every ~20 minutes) and cover a broad range of sizes and amplitudes. They are no hotter than their surroundings and they are Alfvenic in nature, with correlated changes in the magnetic field and velocity.

These results are qualitatively similar to recent simulations of the distant signatures of coronal jets [4,5: see Figure 3], where the plasma outflow rapidly merges with the ambient plasma, but the Alfvenic perturbation travels deep into the solar wind. Since the Helios spikes occur around 5% of the time, have an average kinetic energy 35% higher than the ambient wind and carry around 7% of its total kinetic energy, this implies that they are a non-negligible contribution to the total energy budget of solar wind from coronal holes.

Conclusions

The short velocity spikes in the solar wind at 60 solar radii might be the long-sought signatures of coronal jets. If this were the case, it would open up the possibility of directly measuring these outflows and their properties and implications for the overall solar wind flow.

In the very near future, measurements from Parker Solar Probe at 35 solar radii, nearly twice as close as those considered here, should reveal the fine scale structure of these velocity spikes. Combined with remote observations of the coronal source, either from near-Earth telescopes such as SDO or in future the Solar Orbiter mission, these data should provide unambiguous evidence for the origin of these enigmatic structures.

References

• [1] Extended MHD modelling of the steady solar corona and the solar wind, T. I. Gombosi et al., Living Reviews in Solar Physics, doi:10.1007/s41116-018-0014-4, 2018
• [2] Reprocessed Helios core proton data are available at http://helios-data.ssl.berkeley.edu/data/E1_experiment/New_proton_corefit_data_2017/
• [3] Short, large-amplitude speed enhancements in the near-Sun fast solar wind, T. S. Horbury et al., Mon. Not. Roy. Astron. Soc., Vol: 478, Pages: 1980-1986, doi:10.1093/mnras/sty953, 2018
• [4] Reconnection-driven Magnetohydrodynamic Turbulence in a Simulated Coronal-hole Jet, V. M. Uritsky et al., Ap. J., doi:10.3847/1538-4357/aa5cb9, 2017
• [5] A universal model for solar eruptions, P. F. Wyper et al., Nature, doi:10.1038/nature22050, 2017

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Author: Scott Houston, David Jess, Samuel Grant, Krishna Prasad (Queen’s University Belfast) Andrés Asensio Ramos (IAC), Christian Beck (NSO), Aimee Norton (Stanford).

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Introduction

Studies of the dynamic, oscillatory behaviour of magnetic fields in the solar atmosphere are in their relative infancy, with previous analyses focusing on strong photospheric magnetism. The chromosphere adds a layer of complexity to the energy embedded within the magnetic fields, including omnipresent dissipation, which has implications for thermal connectivity to the corona. The introduction of advanced observational instrumentation, in conjunction with newly developed inversion techniques, has enabled studies of the interaction of non-linear shock fronts in sunspot umbrae (umbral flashes) with the local umbral magnetic field [1][2][3].

Chromospheric observations of sunspot umbrae offer an exceptional view of magneto-acoustic shock phenomena and their impact on the surrounding plasma. Umbral flashes (UFs) are ideal candidates to study magnetic field fluctuations in the chromosphere, since they are energetic, highly non-linear, occupy a large area in the umbra and display well-defined properties. UFs exhibit a periodicity of approximately 3 minutes, which is a consequence of their origin in upwardly propagating magneto-acoustic p-mode waves traversing the density stratification of the lower solar atmosphere, subsequently forming shocks.

In this nugget, we present the results from the first large-scale statistical study of vector magnetic field perturbations, arising from UFs in the chromospheric umbra of a sunspot. We employ the HAZEL [4] inversion code to provide insights into the dynamic fluctuations of both longitudinal and transverse components of the vector magnetic field resulting from shock front interactions.

Observations and Analysis

Using the Dunn Solar Telescope, we acquired high cadence images (ROSA), diffraction limited precision spectro-polarimetry of the chromospheric He I 10830Å line (FIRS) and simultaneous spectroscopic imaging of the Ca II 8542Å absorption profile (IBIS) in the leading spot of NOAA 12565. Calibrated and co-aligned images in ROSA 4170Å continuum, IBIS Ca II 8542Å blue-wing and IBIS Ca II 8542Å line-core, displayed in Figure 1, span the low photosphere into the high chromosphere. The location of the FIRS slit is contoured over the umbral centre.

The high temporal and spectral resolution of IBIS allowed for the identification of 298,091 individual umbral flash pixels in the sunspot umbra. These detected IBIS UFs were co-registered in time and position with FIRS, resulting in 12,988 individual spectro-polarimetric flash spectra. The right panel of Figure 1 displays the characteristic ‘sawtooth’ pattern of an UF in velocity space, indicating that UF signatures are prevalent in the high chromospheric signals captured by the FIRS spectra.

Use the HAZEL code to invert the He I 10830Å Stokes profiles, we investigated the effects of UFs on the local environment, deriving the magnetic field strength, the inclination and azimuth of the magnetic field vector, the Doppler velocity of the plasma and the thermal (velocity) broadening of the sampled plasma.

Relating Temperature and Magnetism

Fluctuations in temperature and magnetic field strength (Figure 2) suggest two distinct populations in our data, related to the filling factor of the shocked plasma in each FIRS pixel. These populations are isolated and displayed in the bottom panels of Figure 2. The first population is classified as ‘edge’ pixels, which have a non-flashing pixel in one or both neighbouring pixels. Pixels in the second, identified as ‘central’ pixels, are bounded by positive UF identifications on both sides.

Edge pixels exhibit a positive correlation between temperature and magnetic field fluctuations. The embedded flux tube scenario can be used to explain this outcome, where initially a flux tube only partially fills an edge pixel. The non-linear shock then provides increased pressure, causing expansion of the flux tube, filling more of the pixel with greater magnetism, while at the same time increasing localised plasma temperatures. The linear relationship can be explained through thermodynamic considerations in an atmosphere dominated by magnetic pressure, leading to an estimate of the sunspot’s chromospheric adiabatic index of γ=1.12 ±0.01.

The second population (pixels filled entirely with shocked plasma) shows an anti-correlation between temperature and magnetic field strength. A decrease in field strength, simultaneous with an increase in temperature, is consistent with previous studies [5], and is also consistent with the viewpoint of shock formation in the lower atmosphere [6].

Transverse Magnetic Field Perturbations

By inverting Stokes Q and U spectra we performed the first examination of the transverse components of the magnetic fields during UF events. The right panel of Figure 3 displays the results of the synthesis, with B_x and B_y denoting the magnetic field components in the solar east-west and north-south directions respectively. The strongest magnetic fields (B > 1600G) are associated with the largest magnitude transverse fluctuations, and are preferentially orientated along the north-north-west to south-south-east direction. However, the weaker fields (B < 1600G) are orientated along the east-north-east to west-south-west direction.… continue to the full article

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