# Searching for evidence of subchromospheric magnetic reconnection on the Sun

Baker, D. University College London, Mullard Space Science Laboratory,
Holmbury St. Mary, Dorking, Surrey, RH5 6NT, UK  van Driel-Gesztelyi, L
University College London, Mullard Space Science Laboratory, Holmbury St.
Mary, Dorking, Surrey, RH5 6NT, UK  LESIA, Observatoire de Paris, Université
PSL, CNRS, Sorbonne Université, Univ. Paris Diderot, Sorbonne Paris Cité, 5
place Jules Janssen, 92195 Meudon, France  Konkoly Observatory, Research
Centre for Astronomy and Earth Sciences, Konkoly Thege út 15-17., H-1121,
Budapest, Hungary  James, A. W. University College London, Mullard Space
Science Laboratory, Holmbury St. Mary, Dorking, Surrey, RH5 6NT, UK  Démoulin,
P  LESIA, Observatoire de Paris, Université PSL, CNRS, Sorbonne Université,
Univ. Paris Diderot, Sorbonne Paris Cité, 5 place Jules Janssen, 92195 Meudon,
France  Laboratoire Cogitamus, rue Descartes, 75005 Paris, France  To, A. S. H
ESTEC, European Space Agency, Keplerlaan 1, PO Box 299, NL-2200 AG Noordwijk,
The Netherlands  University College London, Mullard Space Science Laboratory,
Holmbury St. Mary, Dorking, Surrey, RH5 6NT, UK  Murabito, M  INAF Istituto
Nazionale di Astrofisica, Osservatorio Astronomico di Roma, 00078, Monte
Porzio Catone (RM), Italy  Space Science Data Center (SSDC), Agenzia Spaziale
Italiana, via del Politecnico, s.n.c., I-00133, Roma, Italy  Long, D  School
of Physical Sciences, Dublin City University, Glasnevin Campus, Dublin, D09
V209, Ireland  Brooks, D. H  Department of Physics $\&$ Astronomy, George
Mason University, 4400 University Drive, Fairfax, VA 22030, USA  University
College London, Mullard Space Science Laboratory, Holmbury St. Mary, Dorking,
Surrey, RH5 6NT, UK  McKevitt, J. University College London, Mullard Space
Science Laboratory, Holmbury St. Mary, Dorking, Surrey, RH5 6NT, UK
University of Vienna, Institute of Astrophysics, Türkenschanzstrasse 17,
Vienna 1180, Austria  Laming, J. M. Space Science Division, Naval Research
Laboratory, Code 7684, Washington, DC 20375, USA  Green, L. M  University
College London, Mullard Space Science Laboratory, Holmbury St. Mary, Dorking,
Surrey, RH5 6NT, UK  Yardley, S. L. Department of Mathematics, Physics and
Electrical Engineering, Northumbria University, Newcastle Upon Tyne, NE1 8ST,
UK  Department of Meteorology, University of Reading, Reading, UK  Donostia
International Physics Center (DIPC), Paseo Manuel de Lardizabal 4, San
Sebastián, 20018, Spain  Valori, G  Max Planck Institute for Solar System
Research, Justus-von-Liebig-Weg 3, 37077 Göttingen, Germany  Mihailescu, T
University College London, Mullard Space Science Laboratory, Holmbury St.
Mary, Dorking, Surrey, RH5 6NT, UK  Matthews, S. A  University College London,
Mullard Space Science Laboratory, Holmbury St. Mary, Dorking, Surrey, RH5 6NT,
UK  Kuniyoshi, H  Department of Earth and Planetary Science, The University of
Tokyo, 7-3-1 Hongo, Bunkyo-ku, Tokyo 113-0033, Japan

###### Abstract

Within the coronae of stars, abundances of those elements with low first
ionization potential (FIP) often differ from their photospheric values. The
coronae of the Sun and solar-type stars mostly show enhancements of low-FIP
elements (the FIP effect) while more active stars such as M dwarfs have
coronae generally characterized by the inverse-FIP (I-FIP) effect. Highly
localized regions of I-FIP effect solar plasma have been observed by
Hinode/EIS in a number of highly complex active regions, usually around strong
light bridges of the umbrae of coalescing/merging sunspots. These observations
can be interpreted in the context of the ponderomotive force fractionation
model which predicts that plasma with I-FIP effect composition is created by
the refraction of waves coming from below the plasma fractionation region in
the chromosphere. A plausible source of these waves is thought to be
reconnection in the (high-plasma $\beta$) subchromospheric magnetic field. In
this study, we use the 3D visualization technique of Chintzoglou & Zhang
(2013) combined with observations of localized I-FIP effect in the corona of
AR 11504 to identify potential sites of such reconnection and its possible
consequences in the solar atmosphere. We found subtle signatures of episodic
heating and reconnection outflows in the expected places, in between magnetic
flux tubes forming a light bridge, within the photosphere of the active
region. Furthermore, on either side of the light bridge, we observed small
antiparallel horizontal magnetic field components supporting the possibility
of reconnection occuring where we observe I-FIP plasma. When taken together
with the I-FIP effect observations, these subtle signatures provide a
compelling case for indirect observational evidence of reconnection below the
fractionation layer of the chromosphere, however, direct evidence remains
elusive.

Abundances — Flares

##  1 Introduction

The elemental composition of the Sun’s corona, like other solar-type stars, is
dominated by the First Ionisation Potential (FIP) effect, i.e. in most coronal
structures the abundances of elements with FIP $\leq$10 eV are enhanced
relative to their photospheric abundances. Although there is a large variation
(Mihailescu et al., 2022), active regions (ARs) at or just after their
magnetic flux peak show the highest enhancements (FIP bias) whereas open-field
coronal holes show FIP bias close to 1, i.e. they exhibit similar chemical
composition in the photoshpere and in the corona.

It came as a great surprise that inverse-FIP (I-FIP) composition was detected
in the solar corona in patches over (or in the close vicinity of) sunspot
umbrae in the late phase of flares (Doschek et al., 2015; Doschek & Warren,
2016, 2017; Baker et al., 2019, 2020). I-FIP means depleted/increased
abundance of low-FIP/high-FIP elements, which is common in the coronae of
active cool stars like M dwarfs (e.g. Laming, 2015, and references therein).

In the solar atmosphere, it was shown that the I-FIP effect is due to
depletion of low-FIP elements (Brooks, 2018), in agreement with the
ponderomotive force fractionation model. The model explains the generation of
I-FIP plasma by fast-mode waves propagating upward from the photosphere and
refracting/reflecting at the large density gradient in the chromosphere, where
low-FIP elements are mostly ionised (Laming, 2015, 2021). Baker et al. (2019,
2020) analysed the magnetic evolution of ARs in which I-FIP composition was
observed and found that most of the I-FIP patches are located above forming
(merging) umbrae. This plasma is then evaporated into the corona when the
flare ribbons cross the umbrae. They proposed that such merging umbrae are
likely to undergo subsurface magnetic reconnection as the emerging magnetic
threads are being pushed together by buoyancy. The acoustic waves created
during the reconnection process will mode convert at the plasma $\beta$=1
layer to fast-mode waves (e.g. Schunker & Cally, 2006; Cally, 2007) that will
deplete low-FIP ions from the chromosphere at their reflection location
(Laming, 2021).

Magnetic reconnection in the corona - where the majority of our observations
originate - is taking place in a low-$\beta$ environment. Although a direct
observation of reconnection remains elusive, there is plenty of indirect
evidence from multi-wavelength observations as well as from theory and
simulations that it is taking place in the solar corona and solar wind and it
drives most activity phenomena we observe there (e.g. Benz, 2017; Pontin &
Priest, 2022; Russell, 2023).

The subsurface magnetic reconnection invoked by Baker et al. (2019, 2020) is
assumed to take place below the optical depth $\tau=1$ layer, therefore it has
hardly any observational evidence, not even indirect, besides the I-FIP
coronal composition above merging umbrae (Pontin & Priest, 2022). Analytical
computations (Parker, 1978, 1979) and simulations clearly show, however, that
magnetic reconnection is perfectly possible between flux tubes/ropes below the
photosphere (e.g. Linton, 2006; Murray & Hood, 2007; Tortosa-Andreu & Moreno-
Insertis, 2009; Prior & MacTaggart, 2016; Toriumi & Takasao, 2017; Toriumi &
Wang, 2019). The plasma $\beta$ changes rapidly with depth and can reach
$\beta\approx 10^{4}$ at 2 Mm depth (Chen et al., 2022). However, the
initiation and effects of reconnection differ from the low-$\beta$ case in
that the field lines must be brought together by plasma motions. The main
differences in effects of energy release as a function of plasma $\beta$ were
investigated by Peter et al. (2019) simulating the origin of UV-bursts which
are thought to originate via high plasma-$\beta$ reconnection in the low solar
atmosphere. They find that the peak temperature $T_{peak}$ and plasma-$\beta$
are connected (roughly) by a power law of power $-0.3$. This means that in
high-$\beta$ plasma, magnetic reconnection results in no significant
enhancement of the temperature because plasma density is too high, and thus
the released energy is distributed over many orders of magnitude more
particles than in the solar corona. High density has the same dampening effect
on outflows from the reconnection region; Tortosa-Andreu & Moreno-Insertis
(2009) found that although outflows did appear in their subsurface
reconnection simulations, they were of comparable speed to that of convective
flows.

A crucial link between subsurface reconnection and the fast-mode waves
necessary to create the I-FIP effect is provided by Kigure et al. (2010).
Their 2.5-dimensional magnetohydrodynamical (MHD) simulations show that
component field reconnection between nearly parallel flux tubes can generate
Alfvén and sound waves in a high plasma–$\beta$ environment. The Alfvén waves
originate at the reconnection site and propagate along the reconnected field
lines. Such waves are mostly incompressible (shear) Alfvén waves and they
carry approximately 30$\%$ of the total magnetic energy produced by
reconnection in high $\beta$–plasma (compared to $\approx 40\%$ in low
$\beta$–plasma). Sound waves move across the magnetic field and can also carry
significant energy. The conditions simulated by Kigure et al. (2010) are
similar to what is expected with the coalescing flux of a forming light bridge
- same polarity flux and high $\beta$–plasma.

In this work, we hunt for evidence of reconnection below the fractionation
layer of the chromosphere in NOAA AR 11504. We analyse the evolution of the AR
using multi-wavelength observations obtained from multiple instruments
covering all layers of the solar atmosphere from the photosphere to the
corona. Hinode’s EUV Imaging Spectrometer (EIS; Culhane et al., 2007)
spectroscopic measurements of the AR’s coronal plasma show the presence of
I-FIP plasma during episodes of repeated flux emergence principally located
above light bridges in both polarities. In the context of the ponderomotive
force model, I-FIP plasma composition provides indirect evidence of subsurface
reconnection (Baker et al., 2019, 2020). We infer the large-scale subsurface
structure of AR 11504 using the 3D visualization technique of Chintzoglou &
Zhang (2013) to identify possible locations of reconnection among the flux
tubes forming the light bridges of the coalescing umbrae. We then search for
other signatures, such as energy release and outflows, of high-$\beta$
reconnection in the solar atmosphere related to the location of the coronal
I-FIP plasma.

Figure 1: SDO/HMI line-of-sight magnetogram (a) and continuum (b), SDO/AIA
1600 Å (c), 304 Å (d), 193 Å (e), and 94 Å (f) at 2012 June 13 14:01 UT.
SDO/AIA 304 Å and 94 Å images are shown in reverse colors. Magnetogram is
saturated at $\pm$200 G. Dashed boxes in (d) show the approximate EIS FOV on
June 13-14 (left) and June 15 (right). The animation of the figure contains
these six panels covering from 2012/06/13 00:01 UT to 2012/06/15 05:56 UT with
a data cadence of 5 min and a total viewing duration of 43 sec.

##  2 AR 11504

AR 11504 transited the solar disk from 2012 June 10 to June 20. Initially, the
bipolar AR contained two small sunspots (leading negative/following positive
polarities). However, during the period of June 11-15, significant flux
emergence of $\approx 7\times 10^{21}$ Mx occurred in between the two spots in
the southeastern part of the AR (James et al., 2017). Distinct fragments of
the emerging flux appeared to rotate around the same-polarity (positive) pre-
existing sunspot on the western side of the AR (see Table 2 and Figure 8 of
James et al., 2020). The episodic and large scale flux emergence drove the
evolution of the AR from a bipolar configuration to a large, highly sheared,
complex sunspot group comprising three main sunspots. Flux emergence ceased on
June 15 at about the time that the AR crossed the solar central meridian when
its total magnetic flux was $\approx 10^{22}$ Mx (see Figure 2 of James et
al., 2017), its corrected whole spot area was 1045 millionths of solar
hemisphere (MSH; http://fenyi.solarobs.epss.hu), and its Hale classification
was $\beta\gamma\delta$ (http://helio.mssl.ucl.ac.uk/helio-
vo/solar_activity/arstats-archive/). Figure 1(a,b) shows the Solar Dynamics
Observatory (SDO; Pesnell et al., 2012) Helioseismic and Magnetic Imager (HMI;
Scherrer et al., 2012) photospheric line-of-sight magnetogram and continuum
images at 14:01 UT on June 13 during the flux emergence phase.

The coronal evolution of AR 11504 featured two eruptive events or coronal mass
ejections (CMEs) with M-class flares on June 13 and 14. The eruptive event on
June 14 produced a gradual solar energetic particle event that was detected by
GOES. Both events were preceded by a series of weak, confined homologous
flares and their associated ribbons during the formation of the pre-eruptive
flux ropes. These confined flares are interpreted to be signatures of coronal
magnetic reconnection associated with the formation of the flux ropes that
later erupt as the CMEs (James et al., 2017, 2020). Figure 2 shows the GOES
soft X-ray curve covering the period 00:00 UT June 13 to 08:00 UT June 15. The
approximate timings of the CMEs are indicated by red arrows in the plot.

Figure 1(c–f) contains the chromospheric and coronal images of the AR and its
surroundings just after the peak of an M1.2 flare and the CME eruption on June
13. The locations of the flare loops spanning the AR are identified in the
coronal/flaring SDO/AIA 193 Å and 94 Å passbands and the ribbons are indicated
in the photospheric 1600 Å passband. The evolution of AR 11504 from 00:00 UT
on June 13 to 06:00 UT on June 15 is shown in the online animation of Figure
1. Each frame of the movie consists of six panels: (a) SDO/HMI magnetogram,
(b) continuum, (c) SDO/AIA 1600 Å (d), reverse-color 304 Å, (e) 193 Å, (f) and
reverse-color 94 Å images with FOV [X, Y] = [500$\arcsec$, 600$\arcsec$]. The
data time cadence of the animation is 5 minutes and the viewing duration is 43
seconds. For a full account of the AR’s evolution, its magnetic field
configuration, and flux rope formation/eruption, see James et al. (2017, 2018,
2020).

Figure 2: GOES X-ray flux overplotted with dashed lines to indicate start/end of Hinode/EIS rasters times (see Figure 3). Arrows indicate approximate times of CMEs (James et al., 2020).  Table 1: Hinode/EIS study information (upper) and composition diagnostics (lower). Log Tmax (K) are from the CHIANTI atomic database (Dere et al., 1997) version 10 (Del Zanna et al., 2021). The FIP values are: Ar = 15.76, Ca = 6.11, Fe = 7.87, S = 10.36 eV. Study ID |  $\\#$404 |  $\\#$420
---|---|---
Study name | Atlas$\\_$60 | HPW021$\\_$
|  | VEL$\\_$120x512v1
FOV (arcsec) | 120$\times$160 | 120$\times$512
Exposure time (s) | 60 | 60
Total raster time (m) | 60 | 60
Slit (arcsec) | 2 | 1
Step size (arcsec) | 2 | 2
Raster run time | June 13, 14 | June 15
Ion | Wavelength (Å) | Log T (K)
Ar XIV (high FIP) | 194.40 | 6.60
Ca XIV (low FIP) | 193.87 | 6.65
Ar XI (high FIP) | 188.82 | 6.30
Fe XIV (low FIP) | 264.79 | 6.35
S XIII (intermediate FIP) | 256.69 | 6.45
S XI (intermediate FIP) | 188.68 | 6.35
Figure 3: Hinode/EIS observations of AR 11504 for 2012 June 13 at 11:05 UT
(top), June 14 at 00:14 UT (middle) and June 15 at 00:51 UT (bottom). Units
are in arcsec. Observations on June 13 and 14 are of the AR’s following
(negative) polarity and then of the leading (positive) polarity on June 15
(see Figure 4). Left to right: Ca XIV 193.87 Å intensity map, Ar XIV 194.40
Å–Ca XIV 193.87 Å and Ar XIV 194.40 Å–S XIII 256.69 Å line intensity ratio
maps, Fe XIV 264.79 Å intensity map, Ar XI 188.82 Å–Fe XIV 264.79 Å and Ar XI
188.82 Å–S XI 188.68 Å line intensity ratio maps. The color table of the
ratios (along the bottom) are saturated at an upper level corresponding to the
safe detection of I-FIP, so I-FIP regions are in yellow. Photospheric
composition is approximately orange and coronal composition is green/blue,
however, these depend on the saturation values used as defined by what is
considered to be I-FIP (see end of Section 3.1 and Appendix). (Black stripes
are missing data).

##  3 Hinode/EIS Observations of I-FIP Plasma

###  3.1 Method

Hinode/EIS obtained 26 raster scans of AR 11504 from 2012 June 13–15, three of
which contained patches of I-FIP plasma. Details of the EIS scans where I-FIP
plasma was observed are provided in the upper section of Table 1. The
start/end times of these scans are overplotted on the GOES soft X-ray curve in
Figure 2. Both EIS studies contain a range of possible emission lines deemed
to be appropriate for composition analysis. The available lines from the full
CCD (study name is Atlas$\\_$60) were compared to those high and low FIP line
pairs with similar emissivity temperature dependencies given in Feldman et al.
(2009). The selection criteria for compositional analysis included identifying
different diagnostics for probing plasma at multiple temperatures ranging from
0.3 MK to 4 MK and for the line pairs to be well separated in FIP eV.
Unfortunately, the cooler lines (<1 MK) recommended by Feldman et al. (2009)
were too weak to be used. There are a number of options for plasma at 1.5–4
MK, including the well-known Si X–S X diagnostic for active regions (e.g.
Brooks & Warren, 2011; Brooks et al., 2015). However, recent studies have
shown that in flaring ARs sulfur often acts like a low-FIP element when used
in FIP bias line pairs (Doschek & Warren, 2016; Laming, 2021; To et al.,
2021). The FIP of sulfur is 10.36 eV which places it on the cutoff threshold
between high- and low-FIP elements. The dependence of sulfur’s behavior on the
level of activity in an AR makes it an excellent ion for interpreting plasma
composition observations. Indeed, the behavior of intermediate FIP element of
S is proving to be a key composition diagnostic in the solar wind (Laming et
al., 2019; Parenti et al., 2021), with SEPs (Brooks & Yardley, 2021), and in
flares on all scales (Doschek & Warren, 2017; To et al., 2021, 2024). Its
behavior provides insight into the height of the plasma fractionation in the
chromosphere. In the upper chromosphere/transition region, S acts like the
noble gases with very high FIP and in the low chromosphere, it behaves as a
low FIP element.

Because of the observed high-level of activity of AR 11504, other high-FIP
elements were considered alongside sulfur. As was the case in previous studies
(e.g. Doschek et al., 2015; Doschek & Warren, 2016; Baker et al., 2019, 2020;
To et al., 2021; Mihailescu et al., 2023), high-FIP Ar XIV 194.40 Å–low-FIP Ca
XIV 193.87 Å is used to probe plasma of $\approx$ 4 MK. The behavior and
caveats of this composition diagnostic are covered extensively in the
literature (e.g. Doschek et al., 2015; Doschek & Warren, 2016, 2017; Baker et
al., 2019, 2020). For $\approx$ 2 MK plasma, high-FIP Ar XI 188.82 Å–low-FIP
Fe XIV 264.79 diagnostic was selected from the recommendations listed in
Feldman et al. (2009). The high-FIP and low-FIP elements have similar
emissivities within defined temperature domains and they are separated in
terms of FIP eV as the noble gas argon has the fourth highest FIP in the
periodic table (FIPs: calcium = 6.11 eV, iron = 7.87 eV, and argon = 15.76
eV). Details of the emission line-pairs used to identify I-FIP plasma in AR
11504 are provided in the lower section of Table 1.

Level 1 HDF5 data files were downloaded from
https://eis.nrl.navy.mil/level1/hdf5/ and calibrated using Warren et al.
(2014). Emission lines were fitted using the EIS Python Analysis Code (EISPAC)
software (Weberg et al., 2023) and the included fit templates. Line intensity
ratios for Ca XIV–Ar XIV, Ar XI–Fe XIV, Ar XIV–S XIII, and Ar XI–S XI were
calculated for each pixel in the EIS scans. Figure 3 shows intensity maps of
the low-FIP elements and the line intensity ratio maps at 11:05 UT on June 13,
00:14 UT on June 14, and 00:51 UT on June 15.

The line intensity ratio maps used in this investigation may be subject to
density and temperature effects. Ideally, these effects would be removed by
using pixel–to–pixel differential emission measure (DEM) temperature
corrections and density measurements to make fully calibrated FIP bias maps
(see Baker et al., 2013; Brooks et al., 2015). Due to the difficulty in
constraining DEM curves at temperatures above those of Ca XIV and Ar XIV (log
T > 6.60 K), fully calibrated FIP bias maps were not a realistic option.
Instead, carefully considered thresholds of what constitutes strong I-FIP
plasma for each line pair were determined for the line intensity ratio maps in
Figure 3. Each of the maps has been saturated at levels that are considered to
be safe estimates of I-FIP plasma taking into account such factors as
variations in density and peak temperature in the theoretical ratios for each
diagnostic, assumed abundances, cross-detector sensitivity differences, and
calibration. For Ar XIV–Ca XIV maps strong I-FIP plasma is a ratio $\geq$ 1.75
and for Ar XI–Fe XIV maps, it is a ratio $\geq$0.23. A detailed account of
what constitutes strong I-FIP plasma for the new diagnostic Ar XI–Fe XIV ratio
as well as sulfur-based line pairs is provided in the Appendix.

Figure 4: SDO/HMI and AIA observations of AR 11504 at 2012 June 13 11:36 UT
(top), 2012 June 14 00:46 UT (middle), and 2012 June 15 01:26 UT (bottom).
Left to right: SDO/HMI line-of-sight magnetogram, SDO/HMI continuum, SDO/AIA
1600 Å, and SDO/AIA reverse-color 94 Å maps. Contours (and arrows) are of
I-FIP from the corresponding Ar XI 188.82 Å/Fe XIV 264.79 Å (pink) and Ar XIV
194.40 Å/Ca XIV 193.87 Å (blue) EIS ratio maps in Figure 3. Dashed pink arrow
indicates a few pixels with I-FIP located very close to the light bridge in
the northern region of the following sunspot (HMI continuum image on June 13
only). Black boxes in the right column show the approximate EIS FOV for each
raster.

###  3.2 Location of I-FIP Plasma

Regions of I-FIP plasma are observed on the eastern side of AR 11504 for June
13 and 14 (two top rows of Figure 3) in the higher temperature ($\approx$ 4
MK) Ar XIV–Ca XIV ratio maps (column 2 of Figure 3). The spatial extent of the
patch on June 13 is $\approx$ 10$\arcsec$$\times$10$\arcsec$ centered at [X,Y]
= [-390$\arcsec$, -290$\arcsec$] and on June 14 is $\approx$
30$\arcsec$$\times$20$\arcsec$ at [X,Y] = [-275$\arcsec$, -270$\arcsec$]. The
lower temperature ($\approx$ 2 MK) Ar XI–Fe XIV ratio maps (column 5 of Figure
3) exhibit a similar behavior though the I-FIP regions are larger compared to
the Ar XIV–Ca XIV ratio maps.

Localized I-FIP plasma is again observed at 2 MK but not at 4 MK at 00:51 UT
on June 15, however, the FOV is covering a different region: the western side
of AR 11504. The plasma has evolved to photospheric composition in the Ar
XIV–Ca XIV ratio map. Interestingly, in the Ar XI–Fe XIV ratio map, I-FIP
plasma appears to extend to the south/southwest along loops rooted in the
light bridge containing the I-FIP patch centered at [X,Y] = [85$\arcsec$,
-295$\arcsec$].

All regions of I-FIP plasma are surrounded by a ring of weak
I-FIP/photospheric plasma embedded in the dominant coronal composition plasma.
The precise extent of the I-FIP plasma and the boundaries where the plasma
transitions to photospheric and then to coronal composition depend on the
definition of I-FIP used. It is important to note that I-FIP plasma may have
occurred at locations and times other than in the three scans presented here,
however, due to a number of factors, EIS did not detect it. Some of these
factors include: the limited FOV of EIS, telemetry constraints, pointing
changes, and the lack of composition diagnostics.

Columns 3 and 6 of Figure 3 show line intensity ratio maps where intermediate-
FIP sulfur ions are used instead of the low-FIP calcium and iron ions. In the
case of the Ar XIV–S XIII diagnostic, sulfur behaves like a low-FIP element on
June 13 and 14. This is the case where I-FIP plasma is detected in the Ar
XIV–Ca XIV maps and throughout the core of the AR (cf. columns 2 and 3 of
Figure 3). Sulfur is less fractionated in terms of both degree and spatial
extent when comparing the lower temperature Ar XI–S XI and the higher
temperature Ar XIV–S XIII ratio maps (columns 3 and 6).

The specific locations of the I-FIP plasma within AR 11504 are shown in Figure
4 which displays SDO/HMI magnetic field, HMI continuum, SDO/AIA 1600 Å, and
reverse-color AIA 94 Å maps timed at the mid-points of the EIS scans in Figure
3. The contours of I-FIP plasma from the EIS scans are overplotted on the
SDO/HMI and AIA maps. Pink/blue contours and arrows correspond to I-FIP
patches in the Ar XI–Fe XIV / Ar XIV–Ca XIV line ratio maps.

On the eastern side of AR 11504, patches of I-FIP plasma are observed in the
negative polarity fragments forming the southern of the two following
sunspots. This is the location of the large scale flux emergence that occurred
from June 11–15. A tiny patch is also found at the western edge of the major
light bridge of the northern sunspot (indicated by the dashed pink arrow in
the top row of Figure 4). In the leading (positive) sunspot, the I-FIP plasma
lies above the strong light bridge separating the large pre-existing spot and
a smaller spot that results from new flux emergence. The above analysis of the
photospheric evolution suggests that the spots may represents large flux tubes
that are actually interacting below the visible surface.

Both sites are associated with the footpoints of the flaring loops spanning
the AR and repeated flare ribbon crossings of the coalescing leading and
following sunspots. An example associated with the CME on June 13 is shown in
Figure 1 where the flare ribbons are cospatial with the observed I-FIP plasma
in the top panels of Figure 3. Although the EIS scan was taken mostly prior to
the CME, the 60-min scan time and the W–E scanning direction mean that the
I-FIP patch on the eastern edge of the raster was taken at about 12 UT on the
13th, at the start of the flare/CME event. Repeated flaring on all scales
means that there are frequent occurrences of flare ribbons located above the
forming sunspots on either side of the AR (See the SDO/AIA 1600 Å panel in the
animation of Figure 1).

###  3.3 Relating Hinode/EIS Observations to the Ponderomotive Force
Fractionation Mechanism

Laming (2015) proposed the ponderomotive force associated with Alfvénic waves
as the fractionation mechanism that separates ions from neutrals in the
chromosphere leading to the well-known FIP effect (see also Laming, 2004,
2009, 2012). The ponderomotive acceleration is proportional to the gradient of
$\delta$E2/B2, where $\delta$E is the wave electric field and $B$ is the
ambient magnetic field. Since the magnetic field strength usually decreases
with height, the upward gradient of $\delta$E2/B2 is generally positive which
produces the FIP effect; a negative gradient results in I-FIP plasma.

In the case of the I-FIP effect, fast mode waves originating at/below the
photosphere undergo total internal reflection and refract back downwards in
the lower chromosphere where the Alfvén speed is increasing with height (i.e.
the wave amplitude increases with decreasing density; Laming, 2021). It is
this total internal reflection of waves that provides the downward pointing
ponderomotive force in the model. The force acts only on ions, thereby
depleting the plasma of the easy to ionize low FIP elements and enhancing the
relative abundance of high FIP elements to create the I-FIP effect. The I-FIP
effect plasma is then transported to the corona by chromospheric evaporation
during flaring.

Whether the waves are generated below the chromosphere, at the photosphere, or
subsurface (meaning subphotospheric) is not critical to the ponderomotive
force model. The most important factor is that once the waves propagate above
the plasma $\beta$ = 1 layer, they reflect/refract quickly back downwards and
generate I-FIP fractionation. In the highly complex chromosphere, it is likely
that different geometries would produce I-FIP effect plasma (e.g. Martínez-
Sykora et al., 2023).

Assuming a sufficiently strong source of fast mode waves from below where
high-FIP and low-FIP ions are fractionated in the chromosphere, the
simulations of Laming (2021) show that the degree of I-FIP fractionation
primarily depends on: 1. how much the magnetic field expands from the
photosphere to the corona (expressed as $B$cor/$B$photo), and 2. the
ionization balance of key elements in the lower chromosphere. For the case of
minimal expansion in the magnetic field ($B$cor/$B$photo = 0.7), low FIP
elements, including Fe and Ca, and intermediate FIP S were depleted low in the
chromosphere just above the plasma $\beta$ = 1 layer set at 350 km (see Figure
3 (e,f) of Laming, 2021). This is the height at which the perpendicularly
propagating fast mode waves were undergoing total internal reflection and
creating the negative ponderomotive acceleration. High FIP Ar was unaffected
by the downward directed ponderomotive force as it is essentially neutral in
the model chromosphere. The Hinode/EIS Ar XIV–Ca XIV and Ar XI–Fe XIV line
intensity ratio maps are in full agreement with the simulations. The low FIP
Ca XIV and Fe XIV are depleted relative to their high FIP Ar ion counterparts
as is evident in Figure 3. Like Ca and Fe, S is also largely depleted at the
higher temperature of $\approx$ 4 MK as shown in the Ar XIV – S XIII line
ratio maps. However, S appears to be only partially depleted relative to the
lower temperature of $\approx$ 2 MK. The fractionation of S depends on
collisional interactions with the background gas of neutral H, as opposed to
ionized H higher up in the chromosphere.

Figure 5: 3D visualization of the subsurface magnetic structure of AR 11504
constructed from stacking contours at 1500 G (a) and 2500 G (b) from the Br
component of SHARP vector magnetograms using the method of Chintzoglou & Zhang
(2013). Time runs from 17:00 UT June 9 (top) to 13:00 UT June 20 (bottom) of
each image. Inset in (b) shows the rotating flux tube around the sunspot core
of the positive polarity from a different viewing angle. Key locations for
subsurface reconnection to take place are: 1. flux coalescence in the
following (blue) polarity and 2. at the strong light bridge in the leading
polarity (positive, see Section 4). The animation of the figure shows the
orbital view of the 3D structure first at 1500 G then at 2500 G. The total
viewing time is 48 sec.

Figure 6: Photospheric observations. HMI LOS magnetogram saturated at $\pm$500
G, (a), continuum intensity (b), Dopplergram (c) and zoomed Dopplergram of
boxed region (d).The color scales are in Gauss for the magnetic field and in
km s-1 for the velocity. In panel (d), the black arrow indicates strong
upflows at the light bridge. The two lower rows show SDO/HMI continuum,
running difference (RD) continuum, RD AIA 1700 Å and RD 1600 Å images at 23:59
UT on June 13 (e-h) and 19:57 UT on June 14 (i-l). Red arrows point out the
ridges of brightenings along the light bridge formed by the coalescence of the
leading polarity sunspot. These ridges are not evident in the
cospatial/cotemporal SDO/AIA images of the upper photosphere and transition
region (panels g,h,k,l). The cadence of the running difference movie (not
shown), from which these images were taken, is 45 sec. The animation of the
figure contains panels (a,b,c) covering from 2012/06/12 23:58 UT to 2012/06/15
11:46 UT with a data cadence of 12 min and a total viewing duration of 18 sec.

The I-FIP effect plasma observed by EIS in AR 11504 is consistent with the
predictions of the ponderomotive force fractionation model, suggesting that
there is a source of fast-mode waves originating from below the fractionation
region in the chromosphere. In fact, with its magnetic field strength measured
as $>$2700 G (Pal et al., 2018), the source of the fast-mode waves is
plausibly at/below the photosphere of AR 11504 since the plasma $\beta$ = 1
layer is located lower in regions of strong magnetic field (Gary, 2001; Avrett
et al., 2015; Bourdin, 2017). A candidate for the origin of the required fast
mode waves is subchromospheric reconnection. Such waves are spatially
localized in magnetic flux tubes and are more likely to have sufficient wave
amplitude to cause the fractionation, in contrast to much lower amplitude
acoustic waves excited by convective motions that are present everywhere in
the chromosphere (Laming, 2021).

Other models/simulations also invoke the ponderomotive force as the main
mechanism for fractionation in the chromosphere/transition region (e.g.
Réville et al., 2021; Martínez-Sykora et al., 2023). They focus mainly on the
FIP effect rather than the I-FIP effect, however, and on the whole the EIS
observations of AR 11504 do not appear to contradict their results.

##  4 Search for conditions and signatures of subchromospheric reconnection

As described in Section 1, magnetic reconnection is expected to take place
low-down in the solar atmosphere and even below the photosphere in
high-$\beta$ plasma, where non-parallel magnetic fields have to be brought
together by plasma motions. Simulations show that current sheets can form and
reconnection can take place under such conditions, leading to energy release
and outflows (see Section 1). However, due to the high densities, the
resulting increase in temperature is weak, and the outflow speed is of the
order of the typical convective speed, so the reconnection consequences may be
at the limit of what can be detected. Nevertheless, in the following we make
an attempt to search for signatures of heating and photospheric upflows at the
places where flux emergence is bringing magnetic flux tubes together, e.g. in
merging umbrae, and in particular when new flux is approaching and wrapping
around old flux, forming a prominent light bridge, as we observe in the big
leading spot of this AR. Light bridges are interfaces between magnetic flux
tubes, and in case there is a non-zero angle between the same-polarity
magnetic fields being pushed together by ongoing flux emergence, they are
likely locations for reconnection below the fractionation region. There are
light bridges in both the leading and following regions of this AR, and the
observed I-FIP plasma patches are indeed located in their vicinity. However,
the following polarity light bridges in coalescing umbrae are smaller and less
long-lived than in the big leading spot, so we mainly concentrate our analysis
on the latter, where we can clearly follow the formation and evolution of a
long light bridge with patches of I-FIP plasma above it (due to EIS pointing
choices on June 15). In the following, we look for conditions and consequences
of reconnection there.

###  4.1 Condition: Non-parallel Flux Tubes Pushed Together

The subsurface structure of AR 11504 can be inferred from observations using
the image-stacking technique of Chintzoglou & Zhang (2013). An abridged
description of the image-stacking method is as follows. A 3D data cube is
constructed from SDO/HMI Space-weather HMI Active Region Patches (SHARPs;
Bobra et al., 2014) vector magnetograms. The starting time, t0, is at the top
of the data cube and the 2D magnetograms forming the X–Y plane are added at
progressively lower ‘heights’ in the Z-direction. Contours of radial magnetic
field, Br, from each magnetogram are then stacked along the time dimension.
The overall picture from this analysis is an approximation of the sub-
photospheric 3D magnetic flux structure at time t$\\_$0 as it would result
from a purely solid-body (kinematic) emergence. The solid-body assumption
means that the effects of the turbulent convection zone and emergence through
the photosphere are not captured. Caution is therefore required when
interpreting the 3D structure, especially in the case of weak field.
Nevertheless, this technique can help to visualize the main characteristics of
the evolution of an active region For a full account of the method, see
Chintzoglou & Zhang (2013) and references therein.

The visualization of the 3D structure of AR 11504 from 2012 June 9–20 is
displayed in Figure 5(a,b) and the corresponding online animation. The movie
contains an orbital view of the volume containing the respective stacks of Br
magnetic contours overplotted on HMI vector magnetograms on June 20. Red/blue
represent positive/negative magnetic field. It lasts for 48 sec.

Contour levels of 1500 G and 2500 G (Figure 5(a,b)) were selected to highlight
different features of the emerging structure. The Br stacking of 2500 G
contours emphasizes the larger scale magnetic structure, the so-called
‘trunk’, whereas, the stacking of 1500 G contours provides a sense of the
smaller ‘branches’ of the structure (Chintzoglou & Zhang, 2013).

Light bridges formed as individual flux tubes rotated around the core of the
leading sunspot. One of the main light bridges is indicated in the red Br
stack in the upper right panel of Figure 5(b) and corresponds to the light
bridge of the positive polarity evident in the continuum image (Figure 6(b)).
I-FIP is observed in the EIS line ratio maps at 00:51 UT on June 15 precisely
above this light bridge separating the rotating flux tubes (Figure 4, bottom
row). Furthermore, the loops containing I-FIP plasma appear to be rooted very
close to the light bridge (Figure 3, bottom row).

A more detailed view of the topology is provided by the stack of 1500 G
contours in the 3D visualization. The region of large scale flux emergence
mentioned in Section 2 is highlighted in the negative/following polarity
(blue) in Figure 5(a). Before the data gap on June 13, numerous individual
strands of emerging flux are clearly identifiable in the following sunspot(s)
unlike in the more coherent leading one. After the data gap, the flux
emergence has stopped in the south/southeastern section of the AR and the flux
has coalesced to form two coherent (blue, negative polarity) sunspots. This is
the location where I-FIP was observed on June 13 and 14 while flux
emergence/coalescence was taking place (Figure 4, top and middle rows).

The 3D structure in Figure 5(a,b) can be interpreted as an approximation of
the subphotospheric structure that, by emerging rigidly through the
photosphere, would produce the observed photospheric evolution. Within the
limits of such a visualization, one can still infer the complex dynamical
evolution of the coalescing magnetic flux strands that finally result in the
observed AR’s $\beta\gamma\delta$ Hale classification. For instance, the red
magnetic flux tubes that are pushed together to form the leading polarity have
the same, dominant radial component, but component-reconnection can still
occur between the horizontal component in the volume between them.

Figure 7: White light continuum images of the leading positive sunspot on (a)
2012 June 14 at 00:58 UT and (b) 2012 June 15 at 00:58 UT with arrows that
represent the strength and direction of the horizontal magnetic field. In
panel (a), the horizontal field on either side of the light bridge (approx
$x=-147^{\prime\prime}$, $y=-292^{\prime\prime}$) has small antiparallel
components with very weak horizontal field in the center of the light bridge.
In panel (b), the horizontal field is parallel on either side of the light
bridge, but the horizontal field is weaker in the middle of the light bridge
(approx $x=64^{\prime\prime}$, $y=-299^{\prime\prime}$).

Baker et al. (2019, 2020) posited that strands of sheared field coalescing
during the formation of sunspot umbrae suggests the possibility of subsurface
component reconnection. Component reconnection requires a nonzero component of
anti-parallel field. Such a magnetic field configuration appears in the
subsurface structure of leading and following sunspots. Two key areas are (i)
the region of large scale flux emergence in the following polarity, and (ii)
the flux tube rotating around the core of the leading spot that forms a strong
light bridge in between them. I-FIP plasma is observed in the low corona that
is plausibly connected to subsurface magnetic field in these locations. These
light bridges are locations where the conditions are possible for
subphotospheric reconnection to take place. However, it is important to note
that the 3D visualization of the subsurface structure provides a road map of
potential locations but no information as to whether reconnection has taken
place below the photosphere.

Examining the horizontal magnetic field in the leading sunspot, we observe
small antiparallel field components on either side of the light bridge on June
14 (see Figure 7(a)). On June 15, closer in time to the IFIP observation in
the leading spot, the horizontal field direction is parallel on either side of
the light bridge (see Figure 7(b)). However, the horizontal field strength is
significantly weaker in the center of the light bridge than it is on either
side of the structure which is typical in light bridges (see references in
Section 1). In AR 11504, the horizontal field, Btr is $\sim$200 G in the light
bridge versus $\sim$1500 G in the surrounding regions.

It is interesting to note that if the magnetic field we observe in Figure
5(a,b) is comparable to the field destroyed by reconnection, we can estimate
whether there is enough energy for I-FIP fractionation to result from
reconnection. For example, when the magnetic energy at 1500 G is converted
into waves at a plasma density of 1017 cm-3, then the wave amplitude is
approximately 10 km s-1. This level of wave amplitude is of the right order to
produce the I-FIP effect according to the ponderomotive force fractionation
model. Of course, Figure 5(a,b) shows the magnetic field that survives rather
than what is destroyed by reconnection. However, the numbers, such as we can
constrain them, all support the basic idea of reconnection generating I-FIP
fractionation in AR 11504.

###  4.2 Consequences: Reconnection Outflow and Heating

We search for signatures of subchromospheric reconnection outflows in SDO/HMI
(photospheric) Dopplergrams. The online animation of Figure 6(a–c) shows the
formation of the light bridge as the pre-existing positive-polarity leading
spot is approached by newly emerging positive-polarity flux concentrations.
Each frame contains SDO/HMI LOS magnetograms saturated at $\pm$500 G,
continuum images and Dopplergrams showing photospheric upflows/downflows in
blue/red. The FOV [X, Y] = [250$\arcsec$$\times$170$\arcsec$]. The time period
covered is from 2012 June 12 23:59 UT to June 15 05:47 UT at a time cadence of
12 minutes.

Of particular importance for locating reconnection outflows (appearing as
upflows), in the right panel of the movie and Figure 6(c), SDO/HMI
hmi.V$\\_$720s Dopplergrams are shown. They were downloaded from the the Joint
Science Operations Center (JSOC) website (http://jsoc.stanford.edu).
Dopplergram velocities were calibrated taking into account the orbital motion
of SDO/HMI and the radial velocity of the Sun using the HMI ring diagrams
available from JSOC. The mean Doppler velocity of each map is subtracted from
all pixels within the map (Couvidat et al., 2016; Murabito et al., 2021).

The Doppler maps show effects of the Evershed flow in the penumbra of the main
sunspots. We also see downflows over the leading spot’s umbra, with a faint
trace of the light bridge, where the velocities are mainly around zero. In
these Dopplergrams, we are looking for blue-shifted episodic upflow events
along the light bridge as potential tracers of subchromospheric reconnection
outflows. Previous studies (e.g. Leka, 1997; Toriumi et al., 2015a) reported
upflow velocities of a few $0.1$ to 1.5 km s-1 along light bridges, which have
a weak, more horizontal field than the neighboring umbrae. Furthermore,
convective intrusion driving magnetic reconnection at lower altitudes is also
a possibility (Toriumi et al., 2015b). There is only one clear upflow event
seen in the movie along the light bridge in AR 11504. This upflow is shown in
Figure 6(c,d). It is seen for about 2 hours (on June 13 between 19:10–21:22
UT), around the time of the formation of the light bridge, when EIS was
observing the following polarity region of AR 11504. Unfortunately, there is
no co-temporal corroborating I-FIP observation of this potential reconnection
outflow event.

Next, we search for signatures of heating potentially resulting from
subchromospheric reconnection in running-difference (RD) movies (not shown) of
co-aligned SDO/HMI continuum, SDO/AIA 1700 and 1600 Å images, which represent
increasing height from the photosphere. In Figure 6(e–l), we show snapshots of
SDO/HMI continuum and AIA 1700 and 1600 Å running-difference images. (See
Sunpy documentation for making RD images at
https://docs.sunpy.org/en/stable/generated/gallery/index.html). We find that
in RD continuum the light bridge is clearly seen as a black-and-white
patterned sharp feature, in RD at 1700 Å it is still discernible, but in RD
1600 Å it is not evident at all. As the atmospheric layers these images sample
are increasing in height from continuum to 1600 Å, it gives a hint of low-
altitude weak energy release being present along the light bridge in the
photosphere, which then wanes in the lower atmosphere. The light bridge is
particularly prominent at its formation on June 13 and increased brightness
seems to be present in the vicinity of the upflow event observed in the
Dopplergrams in Figure 6(d), however, the heating at the formation height of
1600 Å does not appear to be as concentrated as the blue-shifted upflows
there.

##  5 Discussion

During a period of large-scale flux emergence in the highly complex and active
AR 11504, the I-FIP effect was observed in plasma at $\approx$ 2 MK and
$\approx$ 4 MK along light bridges in the coalescing umbrae of both AR
polarities. Consistent with previous studies (Baker et al., 2019, 2020), the
regions of I-FIP plasma were localized at the footpoints of flare loops in so-
called monster ARs of high complexity (Hale classification of
$\beta\gamma\delta$), strong total magnetic flux (>1022 Mx), and large sunspot
area (>1000 MSH).

Interestingly, on June 15, the I-FIP plasma area also extends at least
50$\arcsec$ along loops rooted in the vicinity of the strong light bridge of
the leading spot (Figure 3, bottom row). These are the first known
observations showing I-FIP plasma at a distance away from flare loop
footpoints. The timing of the Hinode/EIS observation coincided with the peak
of a C 3.0 class flare that occurred during the extended decay phase of an
M1.9 class flare (see Figure 2), making it possible to observe the I-FIP
plasma along the loops. As noted in Baker et al. (2019, 2020), I-FIP plasma is
observed when it is evaporated into flare loops that cross sunspot umbrae,
however, this can not be directly demonstrated on June 15 as EIS FOV does not
fully cover the footpoints of the loops containing the I-FIP plasma.
Nevertheless, on June 14 at 00:14 UT, the circumstances are similar to those
on June 15. The EIS raster scan covered the negative polarity which hosted
multiple small flaring events during the extended decay phase of an M1.2 class
flare. The I-FIP region (see Figure 3, middle row) is dominated by upflows
characteristic of chromospheric evaporation. Doppler velocities are in the
range = [-100, -20] km s-1 with a mean value of -47 km s-1, much greater than
upflow velocities found in quiescent active regions (Tian et al., 2021). It is
therefore plausible that the pristine I-FIP plasma was evaporated into the
loops during the C3.0 class flare on June 15 and was observed before fully
mixing with the low-FIP enhanced (i.e., high FIP bias) plasma already
contained in the loops. The orange ring of weak I-FIP/photospheric plasma
surrounding the elongated region of I-FIP is an indication of plasma mixing in
the loops (see the Ar XI–Fe XIV line ratio map in the bottom panel of Figure
3). The same feature is notable in all of the line ratio maps shown here as
well as in those of other studies of spatially resolved plasma composition
during flares (e.g. Baker et al., 2019, 2020) where the orange ring encircles
I-FIP patches at loop footpoints (rather than along the loop). In AR 11504,
the I-FIP effect is present in an extended region away from loop footpoint,
suggesting that I-FIP plasma is not restricted to small regions. Though on a
much smaller scale, this is consistent with the findings of Katsuda et al.
(2020) where the I-FIP effect made up the bulk of the plasma in four giant
X-class flares.

Though the presence of I-FIP plasma in the corona is not direct evidence of
subchromospheric reconnection, it nonetheless provides solid clues about the
likely conditions low down in the solar atmosphere/photosphere. I-FIP effect
results when there are sufficiently large fast-mode wave trains generated
below the equipartition layer in order to produce the required downward
ponderomotive force (Laming, 2021). This does not appear to occur under all
flaring conditions in ARs, otherwise, I-FIP plasma would presumably be
observed more frequently over a variety of regions. So the presence of I-FIP
plasma tells us that there are more fast-mode waves generated locally rather
than globally beneath the surface. Subchromospheric reconnection is likely to
be the source of the required wave trains especially in strong sunspots. The
fact that sulfur acts like a low-FIP element in AR 11504 further supports the
argument that the waves are created below the fractionation region in the
chromosphere, and that the fractionation occurs low in the chromosphere. The
few active regions where I-FIP effect plasma has been detected are highly
complex with unusually strong coalescing sunspots. In such strong sunspots,
the plasma $\beta$ = 1 layer is likely to be low down in the photosphere or
even subphotosphere (Carlsson et al., 2019), leaving open the possibility of
reconnection happening below the solar surface. However, reconnection within
light bridges is likely to be in/just above the photosphere. Therefore, highly
localized regions of I-FIP plasma in the corona indicate possible locations of
such low-down reconnection and its consequences in the corona.

In this paper, we were searching for observational evidence –other than I-FIP
plasma observed in the corona, but in tandem with it– of
subchromospheric/high-$\beta$ plasma magnetic reconnection. We found that in
AR 11504 the conditions for subchromospheric reconnection have been met during
the 2nd major episode of flux emergence, between June 13–15, when new flux
approached and merged with pre-existing sunspots and smaller magnetic
concentrations were also merging. In particular, in the leading-polarity spot,
the new flux wrapped around the pre-existing flux (see Figure 5), implying the
possibility of significant non-parallel orientation between the two large
positive-polarity concentrations, likely enabling magnetic reconnection along
their interface. James et al. (2020) found that the rotation rate of the new
flux around the core of the sunspot ($\approx$ 2.6∘ hour-1) and the total
orbit (186∘) peaked during the period when I-FIP plasma was observed in AR
11504. A prominent light bridge signified the interface between the two
magnetic flux tubes/sunspot umbrae, so this was the main target for our search
of signatures of episodic heating and reconnection outflows (potentially
observed as upflows).

Light bridges usually have heating and upflow events (Toriumi et al., 2015b),
but these can be simply attributable to effects of magnetic reconnection
driven by magnetoconvection (Toriumi et al., 2015a). As Tortosa-Andreu &
Moreno-Insertis (2009) has shown, the reconnection outflows are likely to be
comparable in speed with that of the convective motions. We did find an upflow
event lasting for about two hours along the bright newly formed light bridge,
but we did not have co-temporal plasma composition information to corroborate
that it signified subchromospheric/high-$\beta$ plasma reconnection. It is
likely that such subtle effects, upflow and heating evidence, are at the limit
of detectability in HMI observations. From the simulation results of Toriumi
et al. (2015a, b), it is not surprising that we found only weak evidence of
enduring photospheric heating along the light bridge in running-difference
continuum images whose effects diminish with increasing height, and become
unobservable in AIA 1700 and 1600 Å images. Though the observable signatures
of reconnection related to light bridges might be difficult to detect, it is
clear that light bridges are natural locations for reconnection to occur
at/just above the photosphere and thus a potential source of fast mode waves
that can generate I-FIP plasma.

##  6 Conclusions

Direct observational evidence of subchromospheric reconnection remains
elusive, however, a preponderance of indirect evidence has contributed to our
understanding of the conditions required for it to take place and consequences
of the processes associated with it. In AR 11504, there were signs of subtle
heating and weak short-lived reconnection outflows observed in the photosphere
as upflows. We would expect these pieces of evidence of reconnection to be
linked to the presence of the I-FIP effect, although, the observations were
not cotemporal. However, they were all localized in, or at least nearby, light
bridges / in between magnetic flux tubes at the photospheric level.

Among the new directions to pursue is the analysis of changes in sunspot
motions and morphology due to subchromospheric reconnection, employing
techniques such as the 3D visualization of AR subsurface topology used here.
We anticipate that other indications of reconnection in high-$\beta$ plasma
will become more evident with the high resolution observations provided by
DKIST, Solar Orbiter, and SOLAR–C.

Figure 8: Theoretical ratios for ion pairs (a) Ar XIV 194.40 Å–Ca XIV 193.87
Å, (b) Ar XIV 194.40 Å–S XIII 256.69 Å, (c) Ar XI 188.82 Å–Fe XIV 264.79 Å,
and (d) Ar XI 188.82 Å–S XI 188.68 Å corresponding to the line intensity ratio
maps shown in Figure 3. Red lines in (a–d) are coronal and blue lines are
photospheric ratios. Normalized contribution function, G(T, n), convolved with
DEM for I-FIP patches on June 14 at 00:14 UT (middle panels of Figure 3) for
(e) Ar XIV 194.40 Å and Ca XIV 193.87 Å, (f) Ar XIV 194.40 Å and S XIII 256.69
Å (g) Ar XI 188.82 Å and Fe XIV 264.79 Å, (h) Ar XI 188.82 Å and S XI 188.68
Å. Ar XI 188.82 Å–Fe XIV 264.79 Å line intensity ratio map (i) and plot of Ar
XI (black) and Fe XIV (red) normalized intensity profiles (j) along the black
line orthogonal to loops containing I-FIP plasma. Note that the ratio map (i)
is in pixels and is zoomed to the region surrounding the I-FIP plasma along
the loops extending to the south-southwest.

##  Appendix A Method to identify I-FIP plasma

As discussed in Section 3.1, line intensity ratios are used to identify I-FIP
plasma in AR 11504. The six emission lines listed in Table 1 were fitted using
the EISPAC fit templates: ar$\\_$14$\\_$194$\\_$396.6c.template.h5,
ca$\\_$14$\\_$193$\\_$874.6c.template.h5,
ar$\\_$11$\\_$188$\\_$806.3c.template.h5,
fe$\\_$14$\\_$264$\\_$787.1c.template.h5,
s$\\_$$\\_$13$\\_$256$\\_$686.1c.template.h5,
s$\\_$$\\_$11$\\_$188$\\_$675.3c.template.h5. Line intensities were obtained
from the fits to calculate the ratios in each pixel. The intensity ratios were
then compared to the theoretical ratios of the contribution functions at
different densities over the temperature range of log T (MK) = [6.2, 6.8].
Theoretical ratios corresponding to the line intensity ratio maps in Figure 3
are plotted in Figure 8(a–d). The contribution functions were calculated using
the CHIANTI atomic database (Dere et al., 1997) version 10 (Del Zanna et al.,
2021). Photospheric abundances are from Scott et al. (2015a, b) and coronal
abundances are from Schmelz et al. (2012).

In order to better understand the density and temperature effects for each
ratio, we determined the density range and peak temperatures within the I-FIP
patches on June 14 at 00:14 UT. This was done for both diagnostics (2 MK and 4
MK). The relevant spectra were averaged within each of the I-FIP patches
before fitting. Log densities based on the Fe XIII 202 Å–Fe XIII 203 Å density
diagnostic line pair are [8, 9] cm-3. The corresponding peak temperatures were
determined by convolving the contribution functions with the differential
emission measure (DEM). See Baker et al. (2019) for details of the method. The
normalized contribution functions, G(T, n), convolved with DEMs are shown in
Figure 8(e-h). The peak temperatures within the I-FIP patches are consistent
for each line pair – log T = 6.3. Based on the measured densities and peak
temperatures, safe levels for I-FIP is set at a line ratio $\geq$ 0.22 for Ar
XI 188.82 Å–Fe XIV 264.79 Å and $\geq$ 1.75 for Ar XIV 194.40 Å–Ca XIV 193.87
Å, 50$\%$ above their respective maximum photospheric theoretical ratios. A
similar approach was taken for Ar XI-S XI and Ar XIV–S XIII line pairs. I-FIP
plasma is considered to be a ratio $\geq$2.82 for Ar XI-S XI and $\geq$0.23
for Ar XIV–S XIII. These values do not account for a number of factors
including: intensity measurement/calibration errors, changes to the
contribution functions due to abundances, and the differences in the EIS
instrument sensitivity across the long- and short-wave detectors (Del Zanna,
2013; Warren et al., 2014).

Further steps were taken to verify the I-FIP plasma observed along loops at
00:51 UT on June 15. The ratios were calculated after background subtraction
for each element to ensure that the signal-to-noise is sufficient. Following
the method of Brooks et al. (2013), cross-loop intensity profiles were
extracted along a number of lines perpendicular to the loop axis. An example
is shown in panel (j). Intensities were then averaged for 5 stacked rows of
pixels along each line. The background-subtracted profile was fitted with a
Gaussian function in order to obtain the integrated area under each curve.
This was done separately for the Ar XI and Fe XIV intensities. The ratio of
the Ar XI and Fe XIV integrated areas exceeded the I-FIP threshold of 0.22 in
all cases, thereby providing an extra level of confidence that the loops do
indeed contain I-FIP plasma.

Hinode is a Japanese mission developed and launched by ISAS/JAXA,
collaborating with NAOJ as a domestic partner, and NASA and STFC (UK) as
international partners. Scientific operation of Hinode is performed by the
Hinode science team organized at ISAS/JAXA. This team mainly consists of
scientists from institutes in the partner countries. Support for the post-
launch operation is provided by JAXA and NAOJ (Japan), STFC (UK), NASA, ESA,
and NSC (Norway). D.B. is funded under Solar Orbiter EUI Operations grant
number ST/X002012/1 and Hinode Ops Continuation 2022-25 grant number
ST/X002063/1. L.v.D.G. acknowledges the Hungarian National Research,
Development and Innovation Office grant OTKA K-131508. A.W.J. acknowledges
funding from the STFC Consolidated Grant ST/W001004/1. A.S.H.T. is supported
by the European Space Agency (ESA) Research Fellowship. MM has been supported
by the ASI-INAF agreement n. 2022-14-HH.0 and by the Italian agreement ASI-
INAF 2021-12-HH.0 ‘Missione Solar-C EUVST-Supporto scientifico di Fase B/C/D’.
DML is grateful to the Science Technology and Facilities Council for the award
of an Ernest Rutherford Fellowship (ST/R003246/1). The work of DHB was
performed under contract to the Naval Research Laboratory and was funded by
the NASA Hinode program. JML was supported by the NASA Heliophysics Supporting
Research Grant NNH22OB102 and by Basic Research Funds of the Office of Naval
Research. S.L.Y. is grateful to the Science Technology and Facilities Council
for the award of an Ernest Rutherford Fellowship (ST/X003787/1). We recognise
the collaborative and open nature of knowledge creation and dissemination,
under the control of the academic community as expressed by Camille Noûs at
http://www.cogitamus.fr/indexen.html.

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