3.1.1.

Chromosphere fast dynamics

The chromosphere is the “interface layer” between the photospheric plasma dominated by turbulent convective motions, and the tenuous corona where most of the structure is determined by the magnetic field. This atmospheric region still presents crucial observational and theoretical challenges. In the chromosphere, compressible waves excited at the photospheric roots of the granular convection, or through p-mode conversion, steepen and form shocks, accounting for large dynamical excursions and very short timescales, of a few minutes at most,^14–16 requiring a fully time-dependent analysis. Understanding the mechanisms that sustain the chromosphere departure from radiative equilibrium will pave the way to understanding the formation of the super-hot corona and the origin of the solar wind. ADAHELI+ provides high-rate, long-duration observations to cover a large interval in the VIS–NIR range. Table 1 reports the observational requirements related to the study of the dynamics of the chromosphere.

Table 1

Chromosphere fast dynamics: observational requirements.

3.1.2.

Waves and heating of the solar upper atmosphere

The temperature at the top of the Sun’s chromosphere ($\sim \mathrm{10,000}\mathrm{K}$) is higher than at the bottom ($\sim 4300\mathrm{K}$). Different mechanisms have been proposed to explain this rise in temperature, from the dissipation of upward-propagating waves,¹⁷ to resistive dissipation of fine-scale electric currents,¹⁸ to magnetic field reconnection.¹⁹ These waves, usually considered evanescent in a nonmagnetic atmosphere, propagate through so-called magnetoacoustic portals that are generated where the magnetic field is significantly inclined. Such conditions are ubiquitous on the Sun, both in active regions (ARs) and at the boundaries of convection cells.²⁰ In addition to these, magneto-hydrodynamic (MHD) waves in magnetic structures can also significantly participate in the energy budget of the upper layers of the Sun’s atmosphere. In particular, small-scale magnetic elements cover a significant fraction of the solar photosphere^21,22 and harbor different kinds of waves that connect different layers of the solar atmosphere, eventually depositing a significant amount of energy in the upper chromosphere (see, e.g., Refs. 23 and 24 for a complete treatment of the topic). Among the many different kinds of MHD waves that small flux tubes can support (compressive and noncompressive), kink waves are probably the most promising due to their ability to travel long distances before being dissipated.^15,16,25,26 Very recently, observations have revealed the propagation of kink waves in small magnetic elements to the solar chromosphere, with velocity of the order of $6\mathrm{km}/\mathrm{s}$.²⁷ However, although these authors have shown that this propagation is highly nonlinear, and thus subject to dissipation, no signature of energy losses was found between the photosphere and the chromosphere. For these reasons, a new fundamental question arises as to which are the main mechanisms responsible for the dissipation of the energy contained in MHD waves and operating at different heights of the solar atmosphere. To this regard, multiheight high spatial resolution spectropolarimetric observations are needed to study in detail the propagation mechanisms of different kinds of MHD waves and to assess the main dissipation mechanisms involved. Table 2 reports the observational requirements set by the investigation of waves in the solar atmosphere.

Table 2

Waves in the solar atmosphere: observational requirements.

3.2.1.

Prominences, reconnection, and magnetic flux emergence; origin and acceleration of the fast solar wind

Prominences are chromospheric features rooted in the photosphere and extending out of either sides of a primary spine. Magnetic reconnection may create plasma upflows in filaments: as reconnection proceeds, the lower reconnected loops submerge below the solar surface, while the upper ones move upward and carry photospheric plasma with them. Presently, chromospheric jet-like events such as spicules/mottles (in the quiet Sun)^28,29 and fibrils (in ARs)³⁰ are investigated both through the latest generation of ground-based instruments,^31,32 together with the space-based observatories Hinode and Solar Dynamics Observatory.^33,34 Hinode observations of high-velocity spicules have recently revived the discussion on the contribution of these phenomena to coronal heating and solar wind generation.^32,34 However, the short observation times typically allocated to this type of observations on space observatories, due to the shared telemetry among different instruments, and the intrinsic limitations of ground-based observations, hamper further advances on this subject. Experimental confirmation of magnetic reconnection models requires high cadence observations of photospheric and chromospheric lines searching for cancelation of magnetic flux and plasma motions along the line of sight. Uninterrupted space observations with long-time surveys of coronal holes in equatorial and polar regions, available through a dedicated mission such as ADAHELI+, will advance our understanding of where and how the fast solar wind is generated and accelerated.³⁵ Table 3 reports the observational requirements associated with the analysis of chromospheric features.

Table 3

Prominences: observational requirements.

Table 4

Quiet Sun magnetism: observational requirements.

Table 5

X-ray polarimetry of solar flares: observational requirements.

Table 6

VIS-NIR telescope main characteristics.

Table 7

ISODY+ narrowband channel parameters.

3.2.2.

Magnetic flux emergence and quiet Sun magnetism

Magnetic flux emergence is a complex process involving a wide range of time and spatial scales: from the large ARs present during solar maxima (with flux content up to ${10}^{23}\mathrm{Mx}$) that host the most violent phenomena associated with energy release (flares, eruptive prominences, CMEs), to small bipolar flux concentrations (ephemeral regions and granular bipoles, with fluxes from ${10}^{16}$ to ${10}^{19}\mathrm{Mx}$), which populate the quiet Sun at all times during the solar cycle. Recent observations of the solar magnetic surface with space- and ground-based instruments have shown how the quiet Sun magnetic fields can no longer be regarded as sheets of unipolar magnetic flux stretching along the boundaries of large convection cells (i.e., supergranules). Ubiquitous transverse magnetic fields (see, e.g., Refs. 3637.–38) and fine intranetwork mixed polarity fields (see, e.g., Refs. 39 and 40) have been discovered using Hinode instruments.⁴¹ Moreover, the launch of the IMaX spectropolarimeter⁴² on-board the Sunrise balloon-borne solar observatory,^43,44 allowed the resolution of individual kilo-Gauss fluxes⁴⁵ and the exploration of vortex flow motions.⁴⁶ This new information suggests that the quiet Sun magnetic fields are not magnetic debris from decaying ARs but have rather to be generated through some other small-scale mechanisms.⁴⁷ Large fields of view [(FOV) supergranular scale], time-extended observations ($>24\mathrm{h}$), high spatial resolution, and fast time cadence are needed to investigate long-lived structures and see how the intermittent magnetic field appears and disappears. Table 4 reports the observational requirements set by the investigation of quiet Sun magnetism.

3.3.1.

X-ray polarimetry of solar flares

Due to energetic events such as solar flares, the Sun is an astrophysical source with an intense emission of x-rays. Their characterization will advance our understanding of the dynamics of the magnetic fields in the ARs of our star. Magnetic reconnection is the cause of the sudden release of energy in flares and it is responsible for the acceleration of particles,^48–50 including the downward beaming and the upward solar wind. The nonthermal HXR emission, dominating at energies $\ge 15\mathrm{keV}$, is generated by electrons slowing-down in the plasma of the solar atmosphere. Particles radiate via bremsstrahlung and heat the ambient plasma. This emission component is expected to be highly polarized,^51–56 with a polarization degree as high as 40% at 20 keV.⁵⁷ Polarization measurements are directly correlated to the study of the particles’ acceleration directivity and therefore to the understanding of the plasma environment. There have been recent attempts to measure the x-ray polarization with the RHESSI satellite and with the Thomson-scattering polarimeter SPR-N on board of the CORONAS-F satellite.^58–60 Both instruments had a small effective area for polarization measurements and were also heavily affected by the background. Furthermore, the RHESSI spectrometer (not specifically designed to operate as a polarimeter⁶¹) has a high energy threshold for polarimetric measurements (100 keV), thus encountering the problem of the fast decrease of the x-ray flare flux. Neither instruments, therefore, reached the sensitivity needed to achieve significant results. Future missions, e.g., solar orbiter, will not address x-ray polarimetry among their science topics. The photoelectric polarimeter XSPO on board ADAHELI+ is a unique opportunity to complement all the other efforts for the understanding of solar flares. This instrument will perform sensitive x-ray polarimetry in the 15- to 35-keV energy band, in a spectral region where the highly polarized HXR emission starts to overwhelm the SXR component. Additional science topics related to fundamental physics can be addressed with the XSPO: a long-term accumulation of angular resolved data is suitable to reveal the presence of an x-ray-emitting region from the solar disc center, produced by the interaction of axions particles with local magnetic field and also the ARs may prove a good probe for searching for the axionic x-rays.^62,63 Table 5 reports the observational requirements set by the investigation of X-ray polarimetry of solar flares.

5.1.1.

Gregorian Telescope

The telescope relies on a classic Gregorian 2 mirrors layout (see Fig. 2), with an intermediate primary focus allocating a tilted and drilled mirror (FM1), whose function is to reject the radiation outside of the FOV. The elliptic hole of this mirror acts as field stop. FM1 is located in the shadow of M2. A second field stop (FS2) is located behind M2. It stops the light coming directly from the entrance pupil and entering the instrument. A collimator unit (CU) images at infinity the telescope focal plane. Auxiliary components are allocated before the instruments:

Table 6 reports the main characteristics of the telescope.

Fig. 2

Optical layout of the 50-cm Gregory telescope of ISODY+. M1, main mirror; M2, secondary mirror; FM1, folding mirror; SM1, first spherical collimating mirror; CP1, first correcting plate; CP2, second correcting plate; and SM2, aspherical mirror. Overall instrument length is 1580 mm.

Fig. 3

ISODY+ focal plane layout. The focal plane suite consists of a correlation tracker channel (CRTC), the broadband channel (BBC), and the narrowband channel (NBC). SM, spherical mirror; PMU, polarization modulation unit; TTM, tip-tilt mirror; dichroic, dichroic mirror to split the NBC from the BBC; F-P, CSE–FPI; CP1 and CP2, corrector plates; CM, aspherical mirror; FM1 and FM2, folding mirrors; CMOS1, NBC CMOS sensor; BS, beam splitter; CMOS2, BBC CMOS sensor; and CMOS3, CRTC CMOS sensor.

The telescope effective focal length has been chosen as a trade-off between the overall volume, the relief of the instrumental box, and the accommodation inside of focal plane suite.

5.1.2.

Focal plane instruments

The focal plane suite consists of a correlation tracker channel (CRTC), the broadband channel (BBC), and the narrowband channel (NBC). The three channels are arranged in order to optimize the space inside the instrumental box ($600\times 600\times 300{\mathrm{mm}}^3$), located behind the telescope. The channels have been developed taking advantage of a 3-D disposition (see Fig. 3). After the TTM, a suitable dichroic mirror splits the NBC from the BBC. The NBC is able to discriminate very narrow wavebands ($3.8\mathrm{pm}\pm 0.5\mathrm{pm}$ at 810 nm) and retrieve information on the polarization state of the selected photospheric and chromospheric VIS–NIR lines. This is accomplished by using 2 tunable CSE–FPI and a filter wheel (not shown in Fig. 3). The FPIs combine high-spectral resolution with short exposure times and a large FOV, as well as the ability to work in polarized light. After the FPIs, the first corrector plate (CP1) reflects the beam toward the second one (CP2). The aspherical mirror (CM) focuses on CMOS1 the beam after the folding mirror (FM2).

In the BBC, the mirror (FM1) folds the beam toward 2 corrector plates (CP1 and CP2, same as in NBC). Then a beam splitter (BS) divides the incoming light between the proper BBC and the CRTC. The BBC is focused by an aspherical mirror (CM) and a folding mirror (arranged as for the NBC) on CMOS2, with a filter wheel located between BS and CM.

The CRTC is a specular copy of the BBC, with a single interference filter in the place of the filter wheel, whose image is focused on CMOS3.

The main parameters of ISODY+ are summarized in Table 8. Further details on the BBC and the CTCR are given in Refs. 2 and 75.

Table 8

Main parameters on ISODY+ and XSPO instruments.

Table 9

Technology readiness level—platform.

Table 10

Technology readiness level—payload—ISODY+ optical assembly.

Table 11

Technology readiness level—payload—XSPO assembly.

5.2.1.

Gas pixel detector

The GPD is a gas-filled detector devoted to study the polarization of x-ray radiation,^12,13 and it was developed by the Pisa INFN section in collaboration with the INAF-Institute for Space Astrophysics and Planetology. It exploits the dependence of the photoelectric cross section to the polarization of photons to perform the polarimetric measurement. For each x-ray photon absorbed in the gas cell, a photoelectron is emitted with higher probability along the photon polarization direction. The ionization track is drifted and amplified by the gas electron multiplier and eventually collected on a fine subdivided pixel plane ($50\mu \mathrm{m}$ of pitch). The detector has an active area of $1.5\times 1.5{\mathrm{cm}}^2$ and has self-trigger capability that allows for downloading the collected charge only from the small pixel region interested by a track. In Fig. 4 a photoelectron track imaged by the GPD is shown. The analysis of the statistical momenta of the charge distribution allows for calculating the position, projected on the pixel plane, of the absorption point (opposite to the Bragg peak at the end of the track, where a photoelectron loses the largest fraction of its energy) and of the emission direction. In more detail, the image of the track, obtained by means of charge collection of each nonzero pixel of the ASIC subframe, allows for determining the impact point and the emission direction, which is correlated to the polarization of the incoming beam. First, it is calculated the barycenter and the angle that maximizes the second moment with respect to the barycenter (main axis) using all the nonzero pixels. Second, using the sign of the third moment, it is identified which part of the track contains the absorption point which has less charge density with respect to the Bragg peak (end of the track). At this point, the algorithm identifies a region in the initial part of the track using the second moment as a weight, and it evaluates again the barycenter of this part (impact point) and the angle that maximizes the second moment with respect to impact point. This is the retrieved emission direction and its distribution in angle (modulation curve) is related to the polarization of the photon beam,^13,76 which is measured by a ${\cos }^2$ fit to the modulation curve. In the example reported in Fig. 4, the black cross and the continuous black line are the barycenter of the distribution of the charges generated in the GPD by a 22 keV and the main axis of the track, respectively. This information is used to compute the absorption point of the electron on the GPD (marked by a red cross) and the corresponding photoelectron emission direction (drawn as a dashed red line).

Fig. 4

Example of a photoelectron track produced by the absorption of a photon at 22 keV acquired with the GPD filled with an Ar (70%)–DME (30%) gas mixture at the pressure of 2 bar. The upper black cross is the barycenter of the charge distribution and the black solid line is the main axis of the track. The lower red cross is the absorption point and the red dashed line is the emission direction of the photoelectron.

The GPD is, therefore, an instrument capable of performing polarimetric measurements, while acquiring the image of the source, and it has also a moderate spectroscopic capability with an energy resolution of about 20% at 5.9 keV.

The Ar-based gas mixture considered for the ADAHELI+ mission [Ar (60%)–DME (40%) at 3 bar] is suitable to measure the x-ray polarization in the 6- to 35-keV energy range.^77,78 To allow a high quantum efficiency, the absorption gap is 3-cm thick.⁷⁹

5.2.2.

XSPO detector array

XSPO comprises of four photoelectric polarimeters based on the GPD technology. One of them is coupled with a tungsten-coded mask aperture (CMA) needed to localize solar flares on the solar disc with an angular resolution of about 1′. The CMA has a ${70}^{\prime }\times {70}^{\prime }$ fully coded FOV (the solar diameter is about 30′). This detector/mask configuration allows exploiting the GPD imaging capability to localize flares onto the solar disc, but having a minimal impact on the mission profile. The other three GPD units are coupled with simple field angular delimiters (FADs). The solar x-ray flux at energies less than 15 keV is too high for the detector functionality and, moreover, the polarization is expected to be very marginal, therefore a fix multilayer gray filter (as a baseline, $5\mu \mathrm{m}$ $\text{copper foil}+200\mu \mathrm{m}\text{titanium}+200\mu \mathrm{m}\text{aluminum}$) is foreseen in front of each detector beryllium window to remove low-energy photons which otherwise would completely overwhelm the HXR component. Therefore, the polarimeter configuration chosen is effective in the 15- to 35-keV energy band.

XSPO allows reaching a minimum detectable polarization^80,81 of 1% for an X5.1 class flare and 5% for an M5.2 class flare. (Spectral data are taken from Ref. 82 by summing the flux from two footpoints since the polarimeter configuration used is not able to separate spatially their emission.)⁸³ During 2 years of mission operation (from 2017 to 2018) about 20 flares are expected to be observed between the X10 and M5 classes, even if the observations will be performed during the decreasing phase of the solar cycle.

The main parameters of XSPO are summarized in Table 8.