A&A 647, A163 (2021)
https://doi.org/10.1051/0004-6361/202039687
c© ESO 2021
Astronomy
&Astrophysics
VHE gamma-ray detection of FSRQ QSO B1420+326 and modeling
of its enhanced broadband state in 2020
MAGIC Collaboration: V. A. Acciari1, S. Ansoldi2, L. A. Antonelli3, A. Arbet Engels4, M. Artero5, K. Asano6, D. Baack7,
A. Babic´8, A. Baquero9, U. Barres de Almeida10, J. A. Barrio9, J. Becerra González1, W. Bednarek11, L. Bellizzi12,
E. Bernardini13, M. Bernardos14, A. Berti15, J. Besenrieder16, W. Bhattacharyya13, C. Bigongiari3, A. Biland4, O. Blanch5,
G. Bonnoli12, Ž. Bošnjak8, G. Busetto14, R. Carosi17, G. Ceribella16, M. Cerruti18, Y. Chai16, A. Chilingarian19, S. Cikota8,
S. M. Colak5, E. Colombo1, J. L. Contreras9, J. Cortina20, S. Covino3, G. D’Amico16, V. D’Elia3, P. Da Vela17,37, F. Dazzi3,
A. De Angelis14, B. De Lotto2, M. Delfino5,38, J. Delgado5,38, C. Delgado Mendez20, D. Depaoli15, F. Di Pierro15,
L. Di Venere21, E. Do Souto Espiñeira5, D. Dominis Prester22, A. Donini2, D. Dorner23, M. Doro14, D. Elsaesser7,
V. Fallah Ramazani24, A. Fattorini7, G. Ferrara3, L. Foffano14, M. V. Fonseca9, L. Font25, C. Fruck16, S. Fukami6,
R. J. García López1, M. Garczarczyk13, S. Gasparyan26, M. Gaug25, N. Giglietto21, F. Giordano21, P. Gliwny11, N. Godinovic´27,
J. G. Green3, D. Green16, D. Hadasch6, A. Hahn16, L. Heckmann16, J. Herrera1, J. Hoang9, D. Hrupec28, M. Hütten16,
T. Inada6, S. Inoue6, K. Ishio16, Y. Iwamura6, J. Jormanainen24, L. Jouvin5, Y. Kajiwara29, M. Karjalainen1, D. Kerszberg5,
Y. Kobayashi6, H. Kubo29, J. Kushida30, A. Lamastra3, D. Lelas27, F. Leone3, E. Lindfors24, S. Lombardi3, F. Longo2,39,
R. López-Coto14, M. López-Moya9, A. López-Oramas1, S. Loporchio21, B. Machado de Oliveira Fraga10, C. Maggio25,
P. Majumdar31, M. Makariev32, M. Mallamaci14, G. Maneva32, M. Manganaro22, K. Mannheim23, L. Maraschi3, M. Mariotti14,
M. Martínez5, D. Mazin6,40, S. Mender7, S. Mic´anovic´22, D. Miceli2, T. Miener9, M. Minev32, J. M. Miranda12, R. Mirzoyan16,
E. Molina18, A. Moralejo5, D. Morcuende9, V. Moreno25, E. Moretti5, V. Neustroev33, C. Nigro5, K. Nilsson24, D. Ninci5,
K. Nishijima30, K. Noda6, S. Nozaki29,?, Y. Ohtani6, T. Oka29, J. Otero-Santos1, S. Paiano3, M. Palatiello2, D. Paneque16,
R. Paoletti12, J. M. Paredes18, L. Pavletic´22, P. Peñil9, C. Perennes14, M. Persic2,41, P. G. Prada Moroni17, E. Prandini14,
C. Priyadarshi5, I. Puljak27, W. Rhode7, M. Ribó18, J. Rico5, C. Righi3, A. Rugliancich17, L. Saha9, N. Sahakyan26, T. Saito6,
S. Sakurai6, K. Satalecka13, F. G. Saturni3, B. Schleicher23, K. Schmidt7, T. Schweizer16, J. Sitarek11,?, I. Šnidaric´34,
D. Sobczynska11, A. Spolon14, A. Stamerra3, D. Strom16, M. Strzys6, Y. Suda16, T. Suric´34, M. Takahashi6, F. Tavecchio3,
P. Temnikov32, T. Terzic´22, M. Teshima16,42, N. Torres-Albà18, L. Tosti35, S. Truzzi12, A. Tutone3, J. van Scherpenberg16,
G. Vanzo1, M. Vazquez Acosta1, S. Ventura12, V. Verguilov32, C. F. Vigorito15, V. Vitale36, I. Vovk6, M. Will16, and D. Zaric´27,
R. Angioni43,44,?, F. D’Ammando45,?, S. Ciprini43,46, C. C. Cheung47, M. Orienti45, L. Pacciani48, P. Prajapati49, P. Kumar49,
S. Ganesh49, M. Minev50,51, A. Kurtenkov50, A. Marchini52, L. Carrasco53, G. Escobedo53, A. Porras53, E. Recillas53,
A. Lähteenmäki54,55, M. Tornikoski54, M. Berton54,56, J. Tammi54, R. J. C. Vera54,55, S. G. Jorstad57,58, A. P. Marscher57,
Z. R. Weaver57, M. Hart57, M. K. Hallum57, V. M. Larionov58,59,†, G. A. Borman60, T. S. Grishina58, E. N. Kopatskaya58,
E. G. Larionova58, A. A. Nikiforova58,59, D. A. Morozova58, S. S. Savchenko58, Yu. V. Troitskaya58, I. S. Troitsky58,
A. A. Vasilyev58, M. Hodges61, T. Hovatta56,54, S. Kiehlmann62,63, W. Max-Moerbeck64, A. C. S. Readhead61,
R. Reeves65, and T. J. Pearson61
(Affiliations can be found after the references)
Received 15 October 2020 / Accepted 18 December 2020
ABSTRACT
Context. QSO B1420+326 is a blazar classified as a flat-spectrum radio quasar (FSRQ). At the beginning of the year 2020, it was found to be in
an enhanced flux state and an extensive multiwavelength campaign allowed us to trace the evolution of the flare.
Aims. We search for very high-energy (VHE) gamma-ray emission from QSO B1420+326 during this flaring state. We aim to characterize and
model the broadband emission of the source over different phases of the flare.
Methods. The source was observed with a number of instruments in radio, near-infrared, optical (including polarimetry and spectroscopy), ultra-
violet, X-ray, and gamma-ray bands. We use dedicated optical spectroscopy results to estimate the accretion disk and the dust torus luminosity.
We performed spectral energy distribution modeling in the framework of combined synchrotron-self-Compton and external Compton scenario in
which the electron energy distribution is partially determined from acceleration and cooling processes.
Results. During the enhanced state, the flux of both SED components of QSO B1420+326 drastically increased and the peaks were shifted to
higher energies. Follow-up observations with the MAGIC telescopes led to the detection of VHE gamma-ray emission from this source, making
it one of only a handful of FSRQs known in this energy range. Modeling allows us to constrain the evolution of the magnetic field and electron
energy distribution in the emission region. The gamma-ray flare was accompanied by a rotation of the optical polarization vector during a low
-polarization state. Also, a new superluminal radio knot contemporaneously appeared in the radio image of the jet. The optical spectroscopy shows
a prominent FeII bump with flux evolving together with the continuum emission and a MgII line with varying equivalent width.
Key words. gamma rays: galaxies – galaxies: jets – radiation mechanisms: non-thermal – quasars: individual: QSO B1420+326
? Corresponding authors; e-mail: contact.magic@mpp.mpg.de
† Deceased.
Article published by EDP Sciences A163, page 1 of 19
A&A 647, A163 (2021)
1. Introduction
QSO B1420+326, also known as OQ 334, is a blazar located
at redshift of 0.682 (Hewett & Wild 2010). Based on its radio
spectrum it has been classified as a flat-spectrum radio quasar
(FSRQ; Healey et al. 2007). BL Lac objects are divided between
low-, intermediate-, and high-synchrotron peaked (LSP, ISP,
and HSP), while FSRQs are usually only LSP objects. At high
energies (HE; i.e., GeV), FSRQs populate the majority of the
extragalactic gamma-ray sky. Among the associated blazars in
the Fermi-LAT Fourth Source Catalog (4FGL; Abdollahi et al.
2020) there are 694 FSRQs compared to 1131 BL Lac objects.
In the very-high-energy (VHE, &100 GeV) range, despite the
detection of over 60 BL Lac objects by Imaging Atmospheric
Cherenkov Telescopes (IACTs), only about 8 FSRQs are known
to emit in this energy range1. There are a few probable reasons
that contribute to the difference between the number of blazars
detected at HE and VHE. The peak of the gamma-ray emission
in the spectral energy distribution (SED) of FSRQs is usually
shifted to lower energies compared to BL Lac objects (see, e.g.,
Ghisellini 2016). Also, some of the source may have enhanced
internal absorption in the radiation field of the broad line region
(BLR; see e.g., Liu & Bai 2006, but note also Costamante et al.
2018) via the e+e− pair production process. FSRQs are consid-
ered to be more luminous sources than BL Lacs, which means
that they can be detected at greater cosmological distances. How-
ever, for sources located at high redshift (z ∼ 1), the VHE
gamma-ray part of the spectrum is severely absorbed in the pair-
production process on extragalactic background light (EBL; see,
e.g., Domínguez et al. 2011), hampering the discovery potential
in this energy range. Softer gamma-ray spectra resulting from
both effects make detection of FSRQs with IACTs even more
difficult. Finally, FSRQs are known to be extremely variable
(see, e.g., Meyer et al. 2019), which is another complication in
observing these sources with instruments with relatively small
fields of view such as IACTs. The VHE gamma-ray flux has
been seen to vary by as much as two orders of magnitude (see,
e.g., D’Ammando et al. 2019; Zacharias et al. 2019). Emission
variability has been observed down to a timescale of ∼10 min
(Aleksic´ et al. 2011). Due to the strong variability, the cur-
rent most successful approach for studying the VHE gamma-ray
emission of FSRQs is a follow up of alerts of enhanced activity
at lower frequencies. To date, all the cases of discovery of VHE
gamma-ray emission from FSRQs have occurred either during
short flaring activities or longer high states. A notable counterex-
ample is PKS 1510–089, from which persistent VHE gamma-ray
emission was observed during a low flux level at HE (MAGIC
Collaboration 2018). As the number of known VHE FSRQs is
still very small, it is important to observe those objects to look
for common patterns and differences in their emission and to
investigate whether or not the same processes are responsible.
Moreover, the observations need to be multiwavelength, thus
contemporaneously covering the broad energy range of the spec-
trum, which is often difficult to achieve due to fast variability.
QSO B1420+326 is known to be strongly variable, in par-
ticular in the HE range2. A few periods of HE flux enhance-
ment from the source have been observed by Fermi-LAT so
far, the most recent one starting in 2019 December (Ciprini &
Cheung 2020). The high state continued, and based on the HE
flux enhancement, the MAGIC telescopes performed follow-up
1 See http://tevcat.uchicago.edu/, in case of some sources the
classification as FSRQ or BL Lac is however uncertain.
2 https://fermi.gsfc.nasa.gov/ssc/data/access/lat/msl_
lc/source/OQ_334
observations and discovered VHE gamma-ray emission from
QSO B1420+326 in 2020 January (Mirzoyan 2020). The source
is the fourth most distant VHE gamma-ray source known so far.
Thanks to the duration of the enhanced state we were able to
alert other observatories and trace the development of the flare
in a broad range of wavelengths, from radio up to VHE gamma
rays (e.g., Minev et al. 2020; D’Ammando et al. 2020; Fallah
Ramazani et al. 2020). As ejection of new knots in the jets of
FSRQs have often been reported contemporaneously to gamma-
ray flaring activity (see, e.g., Aleksic´ et al. 2014; Lindfors 2015;
Rani et al. 2018), we also organized very-long-baseline radio
interferometry (VLBI) observations of the source following the
high state.
In this paper we report the broadband observations of
QSO B1420+326 triggered by the 2019/2020 high state and
other contemporaneous multiwavelength (MWL) data and their
interpretation. In Sect. 2 we describe the instruments involved in
the campaign, the data taken, and the analysis methods used. The
results of the observations are reported in Sect. 3. In Sect. 4 we
model the broadband emission of the source in different phases
of the high state. Our results are summarized in Sect. 5.
We use cosmological parameters H0 = 67.4 km s−1 Mpc−1,
ΩΛ = 0.6847, and ΩM = 0.315 (Planck Collaboration VI 2020).
2. Observations and data analysis
QSO B1420+326 was observed in 2020 January and February
in a broad energy range by a number of instruments that either
monitor the source, or had responded to a target of opportunity
(ToO) announcement for a high state of the blazar, includ-
ing also publicly released results. We report on the observa-
tions performed in radio (VLBA, OVRO, Metsähovi, Badary
RT32), near-infrared (NIR; CANICA), optical polarization and
photometry (Perkins, LX-200, AZT-8+ST7), optical photome-
try (Siena and Rozhen Observatories, ASAS-SN monitoring,
MIRO/MFOSC-P, REM), optical spectroscopy (LDT), optical
and UV from space satellites (Swift-UVOT and XMM-OM),
X-rays (Swift-XRT and XMM-Newton), HE gamma-rays (Fermi-
LAT), and VHE gamma-rays (MAGIC). To put the results in
the context of earlier measurements we also use archival data
retrieved via the Space Science Data Center3 from the catalogs:
GB6 (Gregory et al. 1996), FIRST (White et al. 1997), NVSS
(Condon et al. 1998), CLASS (Myers et al. 2003), JVASPOL
(Jackson et al. 2007), WISE (Wright et al. 2010), GALEX
(Bianchi et al. 2011), Planck (Planck Collaboration XXVIII
2014), 2RXS (Boller et al. 2016), and SDSS (Albareti et al.
2017). In the archival data sample we also include the lowest
and the highest Swift-XRT states of the source (MJD 57631 and
58831), and the low state observed by Fermi-LAT (integrated
from the mission start, MJD 54683 until MJD 57754).
2.1. MAGIC
MAGIC is a system of two IACTs, each with a mirror dish diam-
eter of 17 m (Aleksic´ et al. 2016a). The telescopes are located on
La Palma in the Canary Islands (28.7◦ N, 17.9◦W), at a height
of 2200 m above sea level.
QSO B1420+326 has been observed by MAGIC on a few
occasions following enhanced states at lower energies. We report
3 https://www.ssdc.asi.it/
A163, page 2 of 19
MAGIC Collaboration: MWL campaign on QSO B1420+326
on observations between 2019 December 31 (MJD = 58848)
and 2020 February 6 (MJD = 58885). The observations con-
sisted of several triggers of the MAGIC Target of Opportu-
nity program and are therefore not continuous, also due to bad
weather periods. The data selection was based on the atmo-
spheric transmission and rates of background events. The total
amount of good-quality data is 14.0 h, out of which 2.9 h, taken
between 2019 December 31 (MJD = 58848) and 2020 January
4 (MJD = 58852), was taken with a special low-energy analog
trigger: SUM-Trigger-II (García et al. 2014). The data are ana-
lyzed using MARS, the standard analysis package of MAGIC
(Zanin et al. 2013; Aleksic´ et al. 2016b). The data selection is
based mainly on the atmospheric transmission measured with
LIDAR (Fruck & Gaug 2015) and on hadronic background rates.
The effect of atmospheric absorption is corrected using LIDAR
information. We analyzed the Sum-Trigger-II part of the dataset
with dedicated low-energy analysis procedures including a spe-
cial image cleaning (the so-called MaTaJu cleaning) with the
cleaning thresholds tuned to the extragalactic field of view of
QSO B1420+326 (Shayduk 2013; Ceribella et al. 2019). For the
part of the dataset during which no signal is detected, we com-
puted upper limits on the flux following Rolke et al. (2005) at
95% confidence level.
2.2. Fermi-LAT
The Large Area Telescope (LAT) is a pair-conversion telescope
launched on 2008 June 11 (MJD = 54628) as one of the two sci-
entific instruments on board the Fermi Gamma-ray Space Tele-
scope (Atwood et al. 2009). Its energy range extends down to
∼30 MeV and up to ∼300 GeV, with peak sensitivity at ∼1 GeV.
In the 4FGL (Abdollahi et al. 2020), QSO B1420+326 is associ-
ated to the gamma-ray source 4FGL J1422.3+3223, which has a
>100 MeV flux of (9.1 ± 1.3) × 10−9 cm−2 s−1 and a power-law
spectrum with photon index 2.38 ± 0.07, which was obtained
from data between 2008 August 4 (MJD = 54682.65) and 2016
August 2 (MJD = 57602.24). QSO B1420+326 does not appear
in any of the Fermi-LAT hard-spectrum source catalogs (see,
e.g., the 3FHL catalog, Ajello et al. 2017), which is consistent
with its relatively steep spectrum.
We use the Python package Fermipy (Wood et al. 2017) to
analyze the Fermi-LAT data. We use Pass8 event data (Atwood
et al. 2013) and select photons of the SOURCE class in a square
region of interest (ROI) of 10◦ × 10◦, centered at the position
of the target source. We perform a binned analysis with ten bins
per decade in energy and 0.1◦ binning in RA and Dec in the
energy range 0.1−300 GeV, adopting the instrument response
functions P8R3_SOURCE_V2. A correction for energy dispersion
is included for all sources in the model except for the isotropic
diffuse components. We apply a cut to include only the gamma-
rays with zenith angle <90◦ to limit contamination from the
Earth’s limb. We include in the model of the region all sources
listed in the 4FGL within 15◦ from the ROI center, along with
the Galactic (Acero et al. 2016) and isotropic diffuse emission
models4.
We first performed a likelihood fit using the full Fermi-
LAT data set available at the time of the analysis, including
events in the time range 2008 August 4 to 2020 February 11
4 gll_iem_v06.fits and iso_P8R3_SOURCE_V2_v1.txt, respec-
tively.
(MJD 54682.66–58890.00). We fit the full spectrum of the tar-
get source, the diffuse models, and the normalization of catalog
sources within 5◦ as free parameters. We also optimized the tar-
get source localization, taking advantage of the ∼3.5 extra years
of data with respect to 4FGL. The detection significance was
estimated with the Test Statistic (TS, Mattox et al. 1996). We
searched for new sources by performing a TS map of the ROI.
We find no significant (TS> 25) new gamma-ray source in this
analysis. Although mild excesses with TS∼ 10 are seen in the
residuals, such fluctuations are to be expected when periods over
long time ranges such as this are considered, and therefore we
choose not to add these excesses as point sources in the ROI
model. The model resulting from this initial fit was used as an
input for the computation of the HE gamma-ray light curves.
The light curves were calculated by performing a likelihood
fit in each time bin. The fitting strategy is designed to adjust the
number of free parameters to the photon statistics available in
each bin. In the Fermi-LAT source catalogs, a source is consid-
ered detected if TS is at least 25. In each light curve bin, we fit
the full spectrum of the target source and the normalizations of
the sources within 3◦ from the ROI center. If the target source
has TS< 25, we progressively restrict the free parameters in the
fit, reloading the average model at each step. First, we reduce
the sources with free normalization to a radius of 1◦, then we
freeze all sources except the target, and finally, if the target is
not significantly detected, we only fit its normalization, leaving
the spectral parameters fixed to the average value from the ini-
tial model. We consider the target source to be detected in a given
time bin if TS> 9, and the signal-to-noise ratio (i.e., flux divided
by its error, or F/∆F) in that bin is larger than two. If this is not
the case, we report a 95% confidence upper limit on the flux.
We report light curves with fixed binning of 1 day and
30 days, and one with adaptive binning (Lott et al. 2012), setting
a constant relative flux uncertainty of 15%. The latter method
provides an estimate of the shortest timescale over which it is
possible to obtain a statistically significant detection and a robust
determination of the spectral parameters of the target.
We calculated 0.1−300 GeV SEDs in the time intervals listed
in Sect. 3 by performing a likelihood fit in several energy bins.
The number of bins is optimized as a trade-off between energy
resolution and photon statistics. We also performed an analy-
sis including data from 2008 August 4 (MJD = 54682) to 2017
January 1 (MJD = 57754), in order to compute a quiescent-
state Fermi-LAT SED to which the high-state ones could be
compared. This time range was chosen based on the monthly
light curve, which shows some early signs of flaring activity
between the second half of 2017 and the first half of 2018 (see
Appendix A.3). This time range is quite similar to the one corre-
sponding to the 4FGL catalog, but provides a small increase in
photon statistics due to the later end time.
Finally, we calculated the probability that each single
gamma ray recorded by the Fermi-LAT is associated with
QSO B1420+326 using the gtsrcprob tool in order to investi-
gate the highest energy photons in the 0.1−300 GeV energy range.
2.3. X-ray
The Neil Gehrels Swift Observatory (Gehrels et al. 2004) car-
ried out 26 observations of the source between 2018 January
20 (MJD = 58138) and 2020 February 10 (MJD = 58889), in
particular on 14 individual days between 2020 January 2 and
February 10 (MJD = 58850–58889). The pointed observations
A163, page 3 of 19
A&A 647, A163 (2021)
were performed with both the X-ray Telescope (XRT; Burrows
et al. 2005, 0.2−10.0 keV) and the Ultraviolet/Optical Telescope
(UVOT; Roming et al. 2005, 170−600 nm). The hard X-ray flux
of this source is below the sensitivity of the BAT instrument
for the short exposures of these observations; therefore, the data
from this instrument are not used. Moreover, the source is not
present in the Swift BAT 105-month catalog (Oh et al. 2018).
All XRT observations were performed in photon counting
(PC) mode. The XRT spectra are generated with the Swift-XRT
data products generator tool at the UK Swift Science Data Cen-
tre5 (for details see Evans et al. 2009). Ancillary response files
are generated with xrtmkarf, and account for different extrac-
tion regions, vignetting, and point-spread function corrections.
We use the spectral redistribution matrices v014 in the Cali-
bration data base maintained by HEASARC. Some of the spectra
have very few photons, and so we are not able to use χ2 statis-
tics. To maintain the homogeneity in our analysis, we grouped
the obtained spectra using the task grppha to have at least one
count per bin and we performed the fit with the Cash statis-
tics (Cash 1979). The data collected during 2019 June 27 and
29 (MJD = 58661 and 58663) were summed in order to have
enough statistics to obtain a good spectral fit. We fit the spectra
with an absorbed power-law using the photoelectric absorption
model tbabs (Wilms et al. 2000), with a neutral hydrogen col-
umn density fixed to its Galactic value (1.08 × 1020 cm−2; Ben
Bekhti et al. 2016). We also applied a log-parabola model to the
XRT data, testing if this model is preferred over a single power-
law model on a statistical basis by applying an F-test. The log-
parabola model is preferred over a simple power-law model only
on 2020 January 28 and at 95% confidence level. However, this
could be due to the low statistics of the single XRT observations.
XMM-Newton (Jansen et al. 2001) observed the source on
2020 January 24 between 04:44:07 and 11:49:07 (MJD 58872.3–
58872.5) for a total duration of 25 ks (observation ID
0850180201). The three EPIC cameras were operated in the
large-window mode with medium filter. The data are reduced
using the XMM-Newton Science Analysis System (SAS v16.0.0),
applying standard event selection and filtering. Inspection of the
background light curves shows that no strong flares were present
during the observation, with good exposure times of 20, 24, and
24 ks for the pn, MOS1, and MOS2, respectively. For each of the
detectors, the source spectrum is extracted from a circular region
of radius 30 arcsec centered on the source, and the background
spectrum from a nearby region of radius 30 arcsec on the same
chip. All the spectra are binned to contain at least 20 counts per
bin to allow for χ2-based spectral fitting. All spectral fits are per-
formed over the 0.3−10 keV energy range using XSPEC v.12.10.1.
The energies of spectral features are quoted in the source rest
frame. All errors are given at the 90% confidence level. The data
from the three EPIC cameras were initially fitted separately, but
because good agreement was found (<5%) we proceeded to fit
them together. Galactic absorption was included in all fits using
the tbabs model.
Three different models were applied: a simple power-law, a
broken power-law, and a log-parabola model. The results of the
fits are presented in Table 1. The F-test shows an improvement of
the fit using both a broken power-law and a log-parabola model
with respect to the simple power-law, with the log-parabola
model providing the best fit.
In order to check for the presence of intrinsic absorption, a
neutral absorber at the redshift of the source was added to this
5 http://www.swift.ac.uk/user_objects
Table 1. Summary of fits to the 0.3−10 keV XMM-Newton spectrum of
the source.
Model Parameter Value
Power-law Γ 1.95 ± 0.01
Flux (0.3–10 keV) (5.2 ± 0.1) × 10−12
χ2/d.o.f. 1627.79/379
Broken power-law Γ1 2.38 ± 0.05
Ebreak 1.22+0.10−0.08
Γ2 1.58 ± 0.04
Flux (0.3–10 keV) (6.0 ± 0.1) × 10−12
χ2/d.o.f. 472.21/377
Log parabola α 2.12 ± 0.01
β −0.60 ± 0.03
Flux (0.3–10 keV) (6.2 ± 0.1) × 10−12
χ2/d.o.f. 432.42/378
Notes. Fits also include absorption fixed at the Galactic value. Flux and
Ebreak (in the source rest frame) are given in units of erg cm−2 s−1 and
keV, respectively.
model, but it does not improve the fit quality and thus is not
required. Moreover, no Fe line was detected in the spectrum. The
90% upper limit on the equivalent width (EW) of a narrow emis-
sion line at 6.4 keV is EW< 10 eV.
2.4. Optical and UV from space-based telescopes
During the Swift pointings, the UVOT instrument observed
QSO B1420+326 in all its optical (v, b, and u) and UV (w1,
m2, and w2) photometric bands (Poole et al. 2008; Breeveld
et al. 2010). For each epoch, possible multiple exposures in the
same filter are first summed with the task uvotimsum and then
analyzed using the uvotsource task included in the HEAsoft
package (v6.28) with the 20201026 release of the Swift/UVOTA
CALDB. We checked if the observations are affected by small-
scale sensitivity problems6. Source counts were extracted from
a circular region of 5 arcsec radius centered on the source,
while background counts were derived from a circular region of
20 arcsec radius in a nearby source-free region.
The Optical Monitor (OM) on board the XMM-Newton satel-
lite observed the source in the u, w1, m2, and w2 filters in
imaging mode together with a fast readout window. The total
exposure times of the imaging observations are: 3500 s (u),
3500 s (w1), 4400 s (m2) and 4400 s (w2). The data were pro-
cessed using the SAS tasks omichain and omfchain. The
UVOT and OM flux densities were corrected for extinction using
the E(B−V) value of 0.010 from Schlafly & Finkbeiner (2011)
and the extinction laws from Cardelli et al. (1989).
2.5. Optical from ground-based telescopes
The REM telescope (Zerbi et al. 2001; Covino et al. 2004),
a robotic telescope located at the ESO Cerro La Silla obser-
vatory (Chile), performed optical photometric observations of
QSO B1420+326 in the period 2020 January 24 to February 6
(MJD = 58872–58885). Observations were carried out with the
Optical Slitless Spectrograph (ROSS2) obtaining three 240 s inte-
gration images in the optical g′, r′, i′ bands. The REM data pre-
sented here were obtained as ToO observations triggered by the
6 https://swift.gsfc.nasa.gov/analysis/uvot_digest/
sss_check.html
A163, page 4 of 19
MAGIC Collaboration: MWL campaign on QSO B1420+326
high gamma-ray flux observed by Fermi-LAT after the MAGIC
detection. Instrumental magnitudes were obtained via aperture
photometry and absolute calibration was performed by means
of secondary standard stars in the field reported by the AAVSO
Photometric All-Sky Survey (APASS) catalog7. Transformation
between theu′ g′ r′ i′ z′ and UBVRI photometric systems was per-
formed using the equations reported in Jester et al. (2005)8.
Optical photometric (BVRI) and polarimetric (R band)
observations were carried out at the 1.83 m Perkins telescope
(Flagstaff, AZ, USA), 40 cm LX-200 telescope (St. Petersburg,
Russia), and 70 cm AZT-8 telescope (Crimea) from 2020 Jan-
uary 23 to April 21 (MJD 58871–58960). The photometric data
were reduced using differential aperture photometry with respect
to comparison stars in the quasar field9. The polarimetric obser-
vations obtained at the Perkins telescope were performed and
reduced in the same manner as described in Jorstad et al. (2010).
The details of polarimetric observations carried out at the LX-
200 and AZT-8 telescopes can be found in Larionov et al. (2008).
Photometric optical observations of QSO B1420+326 were
carried out with the MFOSC-P instrument (Srivastava et al.
2018) mounted on the 1.2 m telescope of Mount Abu IR Obser-
vatory (MIRO10) and used in imaging mode. The source was
observed in B, V , R, and I bands (Johnson-Cousins filters) on
2020 February 2 and 6 (MJD = 58880 and 58885). MIRO is
located at Gurushikhar peak in Mount Abu, India, at an altitude
of 1680 m and is operated by the Physical Research Laboratory
(PRL), Ahmedabad, India.
The optical data from the National Astronomical Observa-
tory (NAO) Rozhen, Bulgaria, were obtained between 2020 Jan-
uary 24 and 26 (MJD = 58872–58874). We used the 2 m RCC
telescope with Andor iKon-L CCD camera (2048× 2048 px,
13.5µm pixel−1) and the 50/70 cm Schmidt telescope with FLI
PL-16803 CCD camera (4096× 4096 px, 9µm pixel−1). Addi-
tional observations were carried out on 2020 January 31 to
February 2 (MJD = 58879–58881) at the Student Astronomi-
cal Observatory (SAO) Plana (Ovcharov et al. 2014) with the
35 cm Newton telescope and SBIG STL-11000M CCD Camera
(4008× 2672 px, 9µm pixel−1). All cameras are equipped with
standard photometric UBVRI Johnson-Cousins filters.
The data were reduced (including bias subtraction, flat-
fielding, and cosmic-ray correction) and analyzed using stan-
dard photometry packages from IRAF11. For each image the
PSF value was measured and aperture photometry was applied.
Standard stars from the SDSS DR12 and VizieR catalogs were
used for photometric calibration after applying transformation
equations12.
The Astronomical Observatory of the University of Siena
observed QSO B1420+326 in its program devoted to optical
photometry of blazars in support of MAGIC. The observatory
runs a remotely operated 30 cm Maksutov-Cassegrain telescope
installed on a Comec 10 micron GM2000-QCI equatorial mount.
The focal plane hosts a Sbig STL-6303 camera equipped with a
3072 × 2048 pixels KAF-6303E sensor; Johnson-Cousins BVRI
filters are available. Multiple 300 s images of QSO B1420+326
were acquired at each visit. After standard dark current subtrac-
7 https://www.aavso.org/apass
8 https://www.sdss.org/dr16/algorithms/
sdssUBVRITransform/
9 See https://vo.astro.spbu.ru/vlar/opt_thumbs/b21420_
1.png
10 https://www.prl.res.in/~miro/
11 http://iraf.noao.edu/
12 http://www.sdss3.org/dr8/algorithms/
sdssUBVRITransform.php#Lupton2005
tion and flat-fielding, images were averaged and aperture pho-
tometry was performed on the average frame by means of the
MaximDL software package. Reference and check stars in the
field of view were selected from the APASS9 (Henden et al.
2016) catalog. The reference R magnitudes were derived from
those reported in the same APASS9 catalog after conversion
between the two different photometric systems, following a for-
mula taken from Munari (2012). Additionally we use publicly
available data in V-band and g-band of ASAS-SN (Shappee et al.
2014; Kochanek et al. 2017).
Conversion of magnitudes to energy fluxes was done using
zero points of Bessell et al. (1998). Correction for the Galac-
tic extinction was applied using the E(B−V) value of 0.011
from Schlafly & Finkbeiner (2011) and the extinction laws from
Cardelli et al. (1989).
We performed observations of optical spectra of the quasar
QSO B1420+326 using the 4.3 m Lowell Discovery Telescope
(LDT; Lowell Obs., Flagstaff, AZ) equipped with the DeVeny
spectrograph and the Large Monolithic Imager (LMI), in
response to the detection of the source at VHE by MAGIC
on 2020 January 21 (MJD = 58869). We employed a grating
setting of 300 grooves per millimeter, which provides spec-
tra from 3300 Å to 7500 Å with a dispersion of 2.17 Å per
pixel, a blaze wavelength of 5000 Å, and a slit width of 2.5′′.
The spectroscopic observations were performed on 2020 Jan-
uary 28 (MJD 58876.578), and February 8 (MJD 58887.497) and
25 (MJD 58904.368). During this month the brightness of the
quasar fell from 14.6 mag to 16.25 mag in R band. Each observa-
tion of the quasar consisted of three exposures of 600 s (900 s on
February 25). Observations of a comparison star HD 126944 of
A type, located ∼86′ from the quasar, were performed before and
after target observations, with two 30 s exposures for each obser-
vation. Bias and flat-field images were obtained regularly. The
LDT allows a switch between different instruments in 2−3 min.
Therefore, photometry of the quasar using the LMI in V and
R filters was performed just before or after spectral observa-
tions to calibrate the spectra. The observations are reduced using
programs written in IDL (v.8.6.1) that implement the technique
described in Vacca et al. (2002) developed for reduction of spec-
tra obtained with SpeX at the NASA Infrared Telescope Facility
on Mauna Kea.
2.6. Near-infrared
The NIR observations were carried out with the camera
CANICA (Carrasco et al. 2017) and the Guillermo Haro 2.1 m
telescope (OAGH) located at Cananea Sonora, Mexico. The
camera is based on a Hawaii detector of 1024 by 1024 pixels
with a plate scale of 0.32 arcsec per pixel. The data are part
of the monitoring program “NIR photometry of AGNs with
Gamma Ray emission detected by Fermi-LAT”. Relative pho-
tometry is obtained with respect to the 2MASS point sources
included in the field of view (5.5 arcmin). Absolute fluxes are
obtained adopting the zero point values of 2MASS derived by
Cohen et al. (2003). The host galaxy is not detected in IR by the
2MASS survey in any of the three bands, resulting in an upper
limit of H ∼ 17.7 mag, much weaker than 13−11 mag observed
by CANICA during the investigated period. Therefore, the effect
of the host galaxy is negligible.
2.7. Radio
The 37 GHz observations were made with the 13.7 m diameter
Metsähovi radio telescope. The detection limit of the telescope
at 37 GHz is on the order of 0.2 Jy under optimal conditions.
A163, page 5 of 19
A&A 647, A163 (2021)
Data points with S/N < 4 are handled as nondetections. The
flux density scale was set by observations of DR 21. Sources
NGC 7027, 3C 274, and 3C 84 were used as secondary calibra-
tors. A detailed description of the data reduction and analysis
is given in Teräsranta et al. (1998). The error estimate in the
flux density includes the contribution from the measurement rms
and the uncertainty of the absolute calibration. We also use pub-
licly available 15 GHz OVRO data (Richards et al. 2011) and
8.63 GHz data from Badary RT32 reported in Kharinov (2020).
We requested Director’s Discretionary Time (DDT) with
the Very Long Baseline Array (VLBA) following the gamma-
ray activity and VHE detection of the quasar QSO B1420+326
and were granted six epochs of observations of the source sep-
arated by approximately 1 month intervals (ID BD227), with
8 h per epoch. Thus far, we have obtained three epochs of
observations under the project, on 2020 March 8, May 10,
and June 6 (MJD = 58916, 58979, 59006). The observations
were performed with all ten antennas in continuum, dual cir-
cular polarization mode at 43 GHz using four intermediate
frequency bands (IFs), each of 64 MHz width. The data were cor-
related at the National Radio Astronomy Observatory (NRAO,
Soccoro, NM) using the VLBA DiFX software correlator. Five
sources were observed at each epoch (QSO B1420+326 3C 279,
3C 345, PKS 1055+18, and B2 1308+326), with 60% of the
time devoted to the main target, QSO B1420+326. The sources
3C 279, 3C 345, and PKS 1055+018 are used for fringe finding
during the correlation. The data were reduced using the Astro-
nomical Image Processing System software (AIPS, van Moorsel
et al. 1996) and Difmap package (Shepherd 1997) in the same
manner as described in Jorstad et al. (2017), except without aver-
aging the final calibrated data over IFs. We used the sources
observed along with QSO B1420+326 to perform absolute cal-
ibration of the electric vector position angle (EVPA), as these
sources are monitored in the VLBA-BU-BLAZAR program13
meaning that their polarization properties at 43 GHz are known.
3. Results
In Fig. 1 we present the MWL light curve summarizing the
evolution of the flare. Based on the VHE state of the source
we define three periods selected for the further spectral analy-
sis: A: 2019 December 29 to 2020 January 5 (MJD = 58846.5–
58853.5): without VHE gamma-ray detection; C: 2020 January
20 to 22 (MJD = 58868.3–58870.3): VHE gamma-ray flare; and
D: 2020 January 25 to February 1 (MJD = 58873.5–58880.5):
detection over longer timescale. In addition we define the fourth
period: B: 2020 January 19 to 20 (MJD = 58867–58868), which
does not have simultaneous MAGIC data, but contains the peak
of the optical and IR flare as well as one of the local peaks of
HE emission. The four periods (referred to throughout the paper
as periods A–D) are summarized in Table 2. In each period we
construct a broadband SED (see Fig. 2).
In the case of gamma-ray data, all observations performed
within a given time window are summed. X-ray data available
from different observations are stacked for the period D, while in
the other periods a single Swift-XRT observation is available and
used. On the other hand, the low-energy data (radio up to UV)
have mostly lower uncertainties and hence are more sensitive to
variability. Therefore, if more than one measurement was taken
at a given time period we average all the measurements and take
the standard deviation of the measurements as the measure of
13 www.bu.edu/blazars/VLBAproject.html
its uncertainty. A similar approach for constructing a broadband
SED was applied by for example MAGIC Collaboration (2018).
3.1. Very high-energy gamma-ray emission
The first detection of the VHE gamma-ray emission
from QSO B1420+326 was achieved on 2020 January 20
(MJD = 58868). A highly significant detection of 14.3σ was
obtained in 1.6 h of effective time (see Fig. 3). In the subsequent
period of 2020 January 26 to February 1 (MJD = 58874–58880)
further hints of signal were obtained, with the highest signif-
icance of the excess (6.6σ) on the night of 2020 January 31
(MJD = 58879), which has also the longest exposure of 2.5 h.
During the flare (period C), the flux observed by MAGIC
above 100 GeV reached (7.8 ± 1.3stat) × 10−11 cm−2 s−1.
The observed spectrum in this period can be described
by a power law: dN/dE = (2.49 ± 0.31stat) ×
10−9(E/100 GeV)−4.22±0.24stat [TeV−1 cm−2 s−1]. Correcting
for the EBL absorption according to Domínguez et al. (2011),
the unattenuated spectrum can be described as dN/dE eτEBL(E) =
(4.04±0.54stat)×10−9(E/100 GeV)−3.57±0.29stat [TeV−1 cm−2 s−1].
After the flare (period D), significant gamma-ray emis-
sion was detected again with 5.6σ, but at about half the flare
level: FD(>100 GeV) = (3.9 ± 0.6stat) × 10−11 cm−2 s−1. The
observed spectrum in this period can be described as: dN/dE =
(0.91±0.13stat)×10−9(E/100 GeV)−3.90±0.25stat [TeV−1 cm−2 s−1].
Correcting for the EBL absorption according to Domínguez
et al. (2011), the spectrum can be described as dN/dE eτEBL(E) =
(1.64±0.22stat)×10−9(E/100 GeV)−2.87±0.36stat [TeV−1 cm−2 s−1].
Despite enhancement of the VHE gamma-ray flux by a factor of
two, the spectral indices during and after the flare are consistent
within 1.5σ. However, it should be noted that in particular in
period D, the uncertainties of the photon index are large.
Before the flare (period A), possibly due to shorter observa-
tions under less favorable zenith angle, no significant emission
was detected and only a 95% C.L. limit of <4.1× 10−11 cm−2 s−1
can be placed for the flux above 100 GeV. The limit is compa-
rable to the emission detected from the source in period D. The
SED of QSO B1420+326 observed by MAGIC in different peri-
ods is shown in Fig. 4.
3.2. High-energy gamma rays
The first detection of a HE outburst from B2 1420+326 was
reported in 2018 December (Ciprini 2018), with a flux increase
by more than two orders of magnitude with respect to the average
4FGL value and significant spectral hardening. A similar spec-
tral hardening was reported in 2019 July (Angioni 2019), where
the first evidence of >10 GeV photons was provided, and again
in 2020 January (Ciprini & Cheung 2020).
The daily Fermi-LAT light curve is shown in Fig. 1, includ-
ing the flux (second panel) and photon index (third panel). Both
the flux and the photon index are significantly variable in this
time interval, based on a simple χ2 test. The Fermi-LAT recorded
a peak daily flux (E > 100 MeV) of (2.6 ± 0.2) × 10−6 cm−2 s−1
on 2020 January 19 (MJD = 58867), corresponding to about
300 times the average value reported in the 4FGL catalog. The
highest-energy photon observed by the Fermi-LAT was recorded
two days prior (2020 January 17, MJD = 58865), with an energy
of ∼150 GeV, providing the first indication of VHE emission
from QSO B1420+32614. Accordingly, the Fermi-LAT recorded
the hardest daily spectrum on the same date, with a photon index
14 The probability of this photon being associated with the target source
is >99.9%, as obtained from the gtsrcprob tool.
A163, page 6 of 19
MAGIC Collaboration: MWL campaign on QSO B1420+326
58845 58850 58855 58860 58865 58870 58875 58880 58885 58890
0.5−
0
0.5
1
1.5
]
-
1
s
-
2
cm
-
10
F [
10
MAGIC >100 GeV
58845 58850 58855 58860 58865 58870 58875 58880 58885 58890
0
10
20
30]-1 s
-
2
 
cm
-
7
F [
10
Fermi-LAT F >0.1GeV
58845 58850 58855 58860 58865 58870 58875 58880 58885 588901.5
2
2.5
Ph
ot
on
 in
de
x
Fermi-LAT index
58845 58850 58855 58860 58865 58870 58875 58880 58885 58890
0
2
4
6
8
]
-
1
 
s
-
2
 
er
g c
m
-
12
F[1
0
0.3-10 keV
Swift-XRT XMM
58845 58850 58855 58860 58865 58870 58875 58880 58885 58890
1.5
2in
de
x
0.3-10 keV index Swift-XRT XMM
58845 58850 58855 58860 58865 58870 58875 58880 58885 58890
0
2
4
6
8
F [
mJ
y] M2 & B-bands UVOT M2 XMM-OM M2 MIRO B UVOT B Perkins B
58845 58850 58855 58860 58865 58870 58875 58880 58885 58890
0
5
10
15
F [
mJ
y] V-band ASAS-SN V&g Rozhen V MIRO VREM V UVOT V Perkins V
58845 58850 58855 58860 58865 58870 58875 58880 58885 58890
0
5
10
F [
mJ
y] R-band
Siena R Rozhen R MIRO R
REM R Perkins++
58845 58850 58855 58860 58865 58870 58875 58880 58885 58890
0
5
10
15
20
Po
lar
iza
tio
n [
%] R-band polarization
58845 58850 58855 58860 58865 58870 58875 58880 58885 58890
150−
100−
50−
0
EV
PA
 [d
eg] R-band EVPA
58845 58850 58855 58860 58865 58870 58875 58880 58885 58890
0
5
10
F [
mJ
y] NIR (I-band)
MIRO I REM I Perkins I
58845 58850 58855 58860 58865 58870 58875 58880 58885 58890
0
20
40
60
F [
mJ
y] NIR (Ks-band)
58845 58850 58855 58860 58865 58870 58875 58880 58885 58890
Time [MJD]
0
1
2F [
Jy] Radio
Metsahovi 37GHz OVRO 15 GHz RT32 8.63GHz
Fig. 1. MWL light curve of QSO B1420+326 between 2019 December 28 (MJD = 58845) and 2020 February 11 (MJD = 58890); see titles and
legends of individual panels. Optical and UV observations are corrected for the Galactic attenuation. The points in red are contemporaneous
(±12 h) with MAGIC observations. The shaded regions show the time ranges of the four considered flare evolution periods (see Table 2). Flux
upper limits in the first two panels are shown with downward triangles.
A163, page 7 of 19
A&A 647, A163 (2021)
Table 2. Periods selected for the SED modeling (the colors of the period
tag correspond to colors used in the figures).
Period MJD Comment
A 58846.5–58853.5 Pre-flare
B 58867–58868 Optical flare
C 58868.3–58870.3 VHE flare
D 58873.5–58880.5 Post-flare
1010 1210 1410 1610 1810 2010 2210 2410 2610
 [Hz]ν
14−10
13−10
12−10
11−10
10−10
9−10
]
-
1
 
s
-
2
 
dN
/d
E 
[er
g c
m
2 E
Archival
MJD 58846.5-53.5 (A)
MJD 58867-8 (B)
MJD 58868.3-70.3 (C)
MJD 58872
MJD 58873.5-80.5 (D)
Fig. 2. MWL SED of QSO B1420+326 in 2020 January and February.
Points follow the colors of the shaded regions in Fig. 1. In addition we
show XMM points taken on 2020 January 24 (MJD = 58872) in black.
MAGIC points are not corrected for the EBL absorption. Gray points
show archival data (most from the ASI Space Science Data Center, but
also including low-state Fermi-LAT spectrum and lowest and highest
X-ray spectrum from Swift-XRT).
of 1.72± 0.08, a flux of (1.2± 0.2)× 10−6 cm−2 s−1 , and a TS of
496. We also note that the photon index was consistently harder
than the 4FGL catalog value (2.38 ± 0.07) during most of the
time range shown in Fig. 1.
The adaptively binned light curve is shown in Fig. 5. The
shortest adaptive bin is centered on 2020 January 19, 00:17:45
(MJD = 58867.012), and has a width of ∼6 h. The highest flux
is recorded in the same bin, reaching a value of (3.6 ± 0.5) ×
10−6 cm−2 s−1, i.e., 400 times higher than the 4FGL value. We
note that a test for spectral curvature was performed in all time
bins of all light curves, and a power-law spectrum was found to
be the best representation in all time intervals.
Additionally, we performed a likelihood fit over the full time
interval included in Fig. 1 to characterize the average source
properties in the HE band in this flaring state. For this time
interval, the LogParabola model is preferred (TScurv = 3615)
with respect to a simple power law to describe the gamma-ray
spectrum of the source. It yields a photon flux (1.09 ± 0.02) ×
10−6 cm−2 s−1, and spectral parameters α = 1.97 ± 0.02, and
β = 0.07 ± 0.0116. We also verified that there are no new signif-
15 TScurv = 2(ln LLP − ln LPL) is used to check whether or not a statisti-
cally significant curvature is detected using a LogParabola model com-
pared with the PowerLaw model, where ln L indicates the logarithm of
the likelihood for each model. A source is considered to have a statisti-
cally significant curvature if TScurv > 25.
16 The functional form of the LogParabola spectral model is dN/dE =
N0(E/Eb)−[α+β log(E/Eb)], where N0 is the normalization, Eb is the refer-
ence energy at which N0 is measured, α is the slope, and β is the curva-
ture parameter.
 ]2 [ deg2θ
0 0.1 0.2 0.3 0.4
e
ve
n
ts
N
0
100
200
300
400
500
Time = 1.63 h
  8.3± = 358.3 
off = 692; NonN
 27.6± = 333.7 exN
σSignificance (Li&Ma) = 14.0
Fig. 3. Distribution of the squared angular distance between the source
nominal and event reconstructed directions (points) and corresponding
background estimation (shaded region) for the night of 2020 January
20 (MJD = 58868) of MAGIC observations. Vertical line shows the cut
value below which the event statistics are given in the inset.
E [GeV]
40 50 60 70 80 90 100 200 300
]
-
1
 
s
-
2
 
dN
/d
E 
[Te
V 
cm
2 E
12−10
11−10
10−10
A (obs.)  
C (obs.)  
C (deabs.)  
D (obs.)  
D (deabs.)  
Fig. 4. SED of QSO B1420+326 observed by MAGIC in periods A,
C, and D (see legend): observed (filled circles) and corrected for EBL
absorption (empty circles).
icant point sources in addition to the initial 4FGL model during
this period.
Finally, as mentioned in Sect. 2, we performed separate like-
lihood fits corresponding to the periods used to build time-
resolved SEDs. The results of these fits are summarized in
Table 3, together with the fits for the quiescent state.
3.3. X-ray
There is no strong variability of the X-ray flux in the inves-
tigated period with an increasing trend from 2020 January 05
(MJD = 58853) to 2020 January 25 (MJD = 58873), where a fac-
tor of two higher peak flux is observed. However, there is clear
variability of the X-ray spectral index from hard values before
the flare to much softer values at the time around the optical and
gamma-ray flares, and then back to hard values. This indicates a
shift of the synchrotron peak of the source, and connected with it
the shift of the crossing point between the synchrotron and the IC
component. This is clearly visible in the X-ray spectra during the
flare and during XMM-Newton observations (see Fig. 2). Results
A163, page 8 of 19
MAGIC Collaboration: MWL campaign on QSO B1420+326
58850 58855 58860 58865 58870 58875 58880
]
-
1
s
-
2
 
cm
-
7
F 
[10
10
20
30
40 F >0.1GeV
58850 58855 58860 58865 58870 58875 58880
Ph
ot
on
 in
de
x
1.5
2
2.5 Photon index
Time [MJD]58850 58855 58860 58865 58870 58875 58880
En
er
gy
 [G
eV
]
10
100 Energies
Fig. 5. Fermi-LAT adaptively binned light curve of QSO B1420+326
in January 2020. The vertical shaded areas follow the color coding of
Fig. 1. Top panel: 0.1−300 GeV photon flux. Middle panel: power-law
photon index (the horizontal shaded area represents the average value
from the 4FGL with its uncertainty). Bottom panel: E > 10 GeV photons
associated with QSO B1420+326 with probability >80%.
Table 3. Results of Fermi-LAT analysis on the different activity states
of QSO B1420+326.
State Flux (a) Photon index TS
Quiescent 0.86 ± 0.13 2.40 ± 0.08 218
A 115 ± 6 2.04 ± 0.04 1942
B 260 ± 20 1.88 ± 0.07 688
C 160 ± 20 1.88 ± 0.06 927
D 121 ± 7 1.93 ± 0.04 2280
Notes. The time periods are defined in Table 2. (a)Total flux in the energy
range 0.1−300 GeV in units of 10−8 cm−2 s−1.
of the spectral fits to individual days of Swift-XRT observations
are given in Appendix A.2.
3.4. Optical
Compared to the historical measurements, the optical emission
is ∼1.5 orders of magnitude higher throughout the investigated
period. Moreover, during that period a strong optical flare is
observed on 2020 January 19 (MJD = 58867) with a variability
timescale of the order of a few days. V-band observations per-
formed during one of the Fermi-LAT peaks show an increase by
nearly an order of magnitude with respect to observations at the
beginning of Period A. Similar variability pattern is seen also in
IR and UV ranges. However, the spectral shape in the IR-UV
range varies during the investigated period, with the spectrum
becoming bluer (harder) during the optical flare (see Fig. 2). This
is consistent with the X-ray behavior of the source that also sug-
gests a shift of the synchrotron peak position. Recently other
occurrences of comparable flaring activity in the optical were
observed (e.g., in July 2019, Marchini et al. 2019), reporting a
R = 13.7 magnitude even slightly brighter than the brightest R
point in Fig. 1.
Fig. 6. Optical spectra of QSO B1420+326 obtained with the LDT; the
spectra are in the frame of the observer.
Also, optical polarization shows interesting behavior with
a dip of the polarization percentage at a few percent level and
a concurrent rotation by ∼150◦. Similar EVPA rotation dur-
ing low-level polarization was also seen contemporaneous with
VHE emission in, for example, PKS 1510–089 (Aleksic´ et al.
2014). High variability of EVPA (down to a timescale of 3 h) was
also seen in PKS 0736+017, another FSRQ, during the period of
VHE gamma-ray detection (H.E.S.S. Collaboration 2020).
3.5. Optical spectroscopy
Figure 6 presents the optical spectra of QSO B1420+326
obtained with LDT at three epochs of different activity states
(taken mostly after the gamma-ray flaring activity). The spec-
tra show the presence of MgII emission line at λ = 4706 Å
(rest wavelength of 2798 Å) and a bump between 3800 Å and
4100 Å (rest frame ∼2260−2450 Å). The latter appears to be part
of an FeII emission complex whose strongest UV lines fall in the
2200−2600 Å range (e.g., Baldwin et al. 2004). The spectra also
include three prominent absorption lines at λ ∼ 4001Å, 4337 Å,
and 4477 Å that intensify as the quasar brightens. The absorp-
tion lines are most likely intervening MgII systems at redshifts
of z ≈ 0.43, 0.55, and 0.60, respectively.
We analyzed the characteristics of the MgII emission line
and FeII bump as a function of continuum brightness. Figure 7
shows an approximate Gaussian fit to the MgII line, while Fig. 8
plots similar modeling of the FeII bump at all three epochs. The
parameters of the Gaussian and velocities of clouds, as well as
the flux of the lines, are given in Table 4, although we are unable
to estimate the velocity of gas producing FeII lines, because the
bump consists of more than 100 lines. There is a significant dif-
ference between the MgII line and FeII bump behavior: (i) The
flux of the MgII line remains constant within 1σ uncertainty
independent of the continuum brightness, while the FeII bump
increases in flux with the continuum level (see Fig. 9); (ii) the
central wavelength of the MgII line fits does not show a shift with
respect to the rest wavelength, while λ◦ of the FeII bump shifts
toward the blue side as the time after the VHE event passes.
Unfortunately, it is not possible to distinguish whether the shift
is due to a relative change of the brightness of lines that form the
FeII bump, or due to gas motion toward the observer; and (iii)
the FWHM of the FeII bump is very stable despite the correla-
tion of its flux with the continuum level, while the velocity of
gas where the MgII line originates increases with time after the
A163, page 9 of 19
A&A 647, A163 (2021)
2740 2760 2780 2800 2820 2840
Wavelength, Α
0.0
0.5
1.0
1.5
F λ
, e
rg
 s-
1 cm
-2
Α
-1
 x 
10
-1
5
MgII
58876
58887
58904
Fig. 7. Observed (solid) and modeled (dotted) MgII line of
QSO B1420+326 in the rest frame at three epochs. The lines on 2020
January 28 (MJD = 58876), February 8 (MJD = 58887), and February
25 (MJD = 58904) are shifted by values of 0.8, 0.3, and 0.1, respec-
tively, for clarity.
VHE event. We also note a significant change of the EW of the
MgII line with the continuum, with EW decreasing as the contin-
uum rises. This questions the identification of QSO B1420+326
as a FSRQ, but Table 4 shows that at the lower levels of activity
EW> 5 Å for the MgII line17.
The increase in the flux of the FeII bump with the continuum
and a possible motion of gas producing FeII lines toward the
observer are quite interesting. The short time-lag between the
continuum and line variability suggests that the FeII emission-
line region is much smaller than the region producing the MgII
line and/or lies close to the line of sight. It is possible that the
nonthermal jet is interacting with an FeII-emitting cloud, while
the MgII emission is excited by the underlying thermal accretion
disk continuum, which varies on a much longer timescale.
We adopt an approach suggested by Ghisellini et al. (2014)
(see also references therein), who used an estimate of the accre-
tion disk (AD) luminosity based on the luminosity of the BLR,
LAD ≈ 10 LBLR. The known flux density of the MgII line, com-
bined with the BLR template constructed by Vanden Berk et al.
(2001) for a composite emission spectrum of a quasar using
SDSS spectra, allow us to estimate the total luminosity of the
BLR in QSO B1420+326 LBLR = (1.8± 0.2)× 1045 erg s−1. This,
for a luminosity distance of 4256.4 Mpc, translates to LAD ≈
2 × 1046 erg s−1, with an uncertainty of a factor of about two, as
suggested by Ghisellini et al. (2014). Interestingly, the obtained
luminosity of the accretion disk is rather high, in the high-end
part of values shown for other sources (Ghisellini et al. 2014).
3.6. Radio
Moderate variability is seen in radio observations (see also
Appendix A.3). OVRO data during the investigated period show
a gradual increase in the flux. No monotonous behavior is seen
in flux at 37 GHz by Metsähovi, but a constant fit can be rejected
at χ2/Nd.o.f. = 51/21. The amplitude of the variability is ∼10%.
The total and polarized intensity VLBA maps of
QSO B1420+326 are presented in Fig. 10. The images
exhibit the bright VLBI core located at the northwestern end of
the jet and a weak extended jet at a position angle of ∼130◦.
17 EW of 5 Å is the classical threshold between BL Lac and FSRQ
objects, see, e.g., Sambruna et al. (1996).
Fig. 8. Observed (black) and modeled (blue) FeII bump of
QSO B1420+326 at three epochs in the rest frame; red lines represent
Gaussian fits to absorption lines, and gray dashed lines indicate the con-
tinuum level.
The total and polarized intensity images are modeled by cir-
cular components with Gaussian brightness distributions. Two
bright features are apparent at each epoch: the core A0, and a
knot K20. We assume that the core is a stationary feature of
the jet and calculate parameters of K20 with respect to the core.
Parameters of the modeling are given in Table 5.
According to Fig. 10 and Table 5, knot K20 is the most
polarized feature of the jet. Figure 11 shows the separa-
tion of K20 from the core as a function of time. Accord-
ing to a linear approximation, K20 moves with a proper
motion µ= 0.30± 0.04 mas yr−1, which translates into superlu-
minal motion down the jet with apparent speed βapp = 12.0± 1.7
in units of c. Such motion suggests that the ejection18 of K20
through the VLBI core occurred on MJD 58831± 21 (2019
December 13). Figure 11 also shows the light curve of K20,
which reveals a very fast decrease of the knot flux density by
a factor of about four in three months, which corresponds to the
timescale of variability of 0.16± 0.03 yr, as defined according
18 Passage of the centroid of K20 through the centroid of A0.
A163, page 10 of 19
MAGIC Collaboration: MWL campaign on QSO B1420+326
Table 4. Parameters of lines.
MJD Line λ◦ Amp FWHM v Flux S cont EW
(1) (2) (3) (4) (5) (6) (7) (8) (9)
58876.578 MgII 2795± 2 0.37± 0.04 18.4± 1.3 1982± 140 13± 3 7.1± 0.2 2± 1
58887.497 MgII 2798± 3 0.36± 0.06 26.0± 2.6 2787± 278 18± 4 3.3± 0.2 6± 3
58904.368 MgII 2797± 2 0.31± 0.02 28.6± 0.7 3074± 75 16± 2 1.5± 0.2 11± 2
58876.578 FeII 2390± 5 0.49± 0.05 91± 3 – 84± 9 9.2± 0.5 –
58887.497 FeII 2357± 5 0.23± 0.06 91± 2 – 58± 5 4.5± 0.2 –
58904.368 FeII 2348± 5 0.15± 0.03 91± 1 – 27± 7 1.8± 0.3 –
Notes. Columns as follows: (1) epoch of observations; (2) emission line; (3) central wavelength of Gaussian fit in Å; (4) amplitude of Gaussian
fit in units of 10−15 erg cm−2 s−1 Å−1; (5) full width of Gaussian at half-maximum (FWHM) in Å; (6) velocity of gas in km s−1; (7) flux of line in
10−15 erg cm−2 s−1; (8) flux density of the continuum at the peak of line in 10−15 erg cm−2 s−1 Å−1; (9) equivalent width in Å.
Fig. 9. Dependence between flux of the FeII bump and flux density of
the continuum at different epochs (points show the measurements, and
the line is a linear fit).
to Burbidge et al. (1974). Applying the formalism proposed by
Jorstad et al. (2005), this timescale of variability and the aver-
age size of K20 (0.084± 0.015 mas according to Table 5) give a
value of the Doppler factor of K20 δ= 33± 9. The latter, along
with the apparent speed of the knot, allow us to estimate the
Lorentz factor, Γb = 19± 9, and viewing angle, Θ◦ = 1.1◦ ± 0.6◦,
of K20. Using the proper motion and size of K20, we can esti-
mate that during the time of VHE emission the upstream edge of
the knot was passing through the centroid of the core. The direc-
tion of polarization in K20 is close to the jet direction, which
suggests that K20 is most likely a moving shock whose sur-
face is oriented transverse to the jet axis. This is supported by
a higher degree of polarization of the knot compared with the
core, which implies ordering of the magnetic field as expected
by such a shock.
4. Spectral energy distribution modeling
FSRQ gamma-ray emission is usually explained in the frame-
work of an external Compton (EC) model. Moreover, the opti-
cal spectroscopy measurements performed during the dimming
phase of the emission (see Fig. 8) show a significant increase
in flux of FeII + FeIII lines. As the source emits up to VHE
range, the most natural target for EC process is the dust torus
(DT) radiation field (see, e.g., Costamante et al. 2018; van den
Berg et al. 2019). DT radiation field, contrary to BLR, would not
strongly absorb the sub-TeV gamma-rays, and thus is the com-
mon assumption for the dominating radiation field in the model-
ing of FSRQs detected at VHE gamma-rays. On the other hand,
the large increase in the optical flux during the gamma-ray flare
can also provide a significant target for the synchrotron-self-
Compton (SSC) process. We therefore investigated a scenario
in which both SSC and EC processes are possible.
Intriguingly, compared to the low-state spectrum, the inves-
tigated synchrotron and gamma-ray peaks are shifted towards
higher energies. Similar behavior has been observed in a
few other FSRQs during enhanced states, e.g., PMN J2345–
1555 (Ghisellini et al. 2013), 4C +49.22 (Cutini et al. 2014),
PKS 1441+25 (Ahnen et al. 2015; Abeysekara et al. 2015),
PKS 1510–089 (D’Ammando et al. 2011), and PKS 0346–27
(Angioni et al. 2019). However, for most of these cases, the peak
frequency did not reach beyond 1014 Hz. Therefore, the behavior
observed in QSO B1420+326, in particular the SED peaking at
a few times 1014 Hz in the optical range during period B, while
not being unique, is still rarely observed in FSRQs. The peak
position traces the electron energy distribution (EED), but it is
also dependent on other physical parameters of the source (e.g.,
on the beaming).
We model the source in a framework of a simple one-zone
model in which a spherical emission region is homogeneously
and isotropically filled with an electron distribution and mag-
netic field. We consider a broken power-law energy distribu-
tion of electrons, i.e., dNe/dγ ∝ γ−p1 for γmin < γ < γb and
dNe/dγ ∝ γ−p2 for γb < γ < γmax. The electrons in the blob are
also exposed to an additional, directional radiation field coming
from the DT. The model assumes ring geometry of the DT, and
thus depends on the distance of the emission region from the
black hole. The SED model was generated with agnpy19 (Nigro
et al. 2020), which implements the synchrotron and Compton
processes following the prescriptions described in Dermer &
Menon (2009), Finke (2016). We fix the Lorentz and Doppler
factors of the blob to Γ = 40 and δ = 40, respectively. We note
that those values are somewhat larger (in particular the Γ fac-
tor of the blob) than the jet parameters estimated using VLBI
observations (see Sect. 3.6). However, the VLBI measurements
are performed a few months after the flaring event and trace
the later evolution of the blob, and therefore some change of
parameters of the jet might have happened in the meantime (in
particular deceleration). The size of the blob is limited by the
variability condition. The values that we use in the modeling,
rb = 3−6 × 1016 cm, correspond to the light-crossing time of
12−24 h, of the order of the timescale of the observed variability.
19 https://github.com/cosimoNigro/agnpy/
A163, page 11 of 19
A&A 647, A163 (2021)
Fig. 10. Total (contours) and polarized (color scale) intensity maps of
QSO B1420+326 obtained with the VLBA at 43 GHz; the peak total
intensity is 1.14 Jy beam−1; the beam is displayed in the bottom left cor-
ner; the contours are 0.25, 0.5, 1, . . . 64, and 90% of the peak total inten-
sity. Linear segments within images show direction of polarization, the
black vertical line indicates the position of the core, A0, and red circles
mark positions of knot K20; 1 mas corresponds to 7.29 pc. We note that
the polarized flux intensities on May 10 and June 6 are multiplied by
factors of three and five, respectively, to reveal locations of peaks of
polarized flux intensity, which are 15 mJy beam−1 and 9.5 mJy beam−1,
respectively, for the May 10 (MJD = 58979) and June 6 (MJD = 59006)
images.
We assume that the emission region is located at a distance of
d = 2.5×1018 cm, that is, of the order of ∼Γrb. We use the accre-
tion disk luminosity Ldisk = 2 × 1046 erg s−1 (see Sect. 3.5) to
estimate the size of the BLR and DT following the scaling rela-
tions of Ghisellini & Tavecchio (2009). We note that while the
optical spectra used in this estimation are not fully simultaneous
with the broadband emission data used for the modeling, the size
of the DT makes the emission quasi-stable at the timescales of
years. The dust torus is simulated as a thin ring with a radius of
RDT = 1.1 × 1019 cm (3.6 pc). As the estimated size of the BLR
is RBLR = 4.4 × 1017 cm (0.14 pc), the emission region is not
affected by the BLR but it is deep in the DT radiation field. We
assume that 0.6 of the disk luminosity is reprocessed in the DT
radiation.
In order to obtain EED in a self-consistent way, we introduce
the acceleration parameter ξ, defined such that acceleration gain
of electrons is (dE/dt)acc = ξcE/RL, where c is the speed of light
and RL is the Larmor radius. The maximum energy of the elec-
trons is obtained from comparing the acceleration energy gain
with energy losses due to IC (in Thomson regime) cooling:
γmax =
√
3ξeB/4σTu′ph, (1)
where B is the co-moving magnetic field in the blob, e is ele-
mentary charge, σT is the Thomson cross-section, and u′ph is the
co-moving energy density of the dominating radiation field. In
the case of the parameter sets used in the modeling, the dominat-
ing radiation field is originating from the DT. Nevertheless, we
explicitly also check a possible limit from the SSC process. The
maximum electron energies are also tested against the dynamical
timescale by comparing the acceleration timescale with the bal-
listic timescale20, Tbal ' rb/c, of the blob crossing its size, and
against synchrotron energy losses. However, neither of the two
processes is dominant in the cases investigated here. As during
Tbal the blob crosses its size, it might escape the region in the
jet (e.g., a stationary shock) where efficient acceleration occurs.
We note that Tbal also determines the timescale of the adiabatic
losses of the electron (following the assumption that the blob
distance is d = rbΓ).
We also introduce a canonical cooling break (steepening of
EED by 1) at the energy where the dominating cooling timescale
is equal to the ballistic timescale, Tbal:
γb = 3mec2/4σTu′phrb, (2)
where me is the electron mass. In order not to overshoot the X-
ray flux with IC photons and still be able to explain the gamma-
ray emission during the periods B, C, and D, it is necessary
to assume that the EED starts at γmin > 1. We apply values
of γmin = 10−15 in the modeling, while for period A we use
γmin = 1. We then tune the magnetic field B and the energy den-
sity of the electrons u′e to reproduce the levels of synchrotron and
IC emission. The index of EED before the break is selected to
roughly reproduce the spectral shape in IR-UV (except of period
D) and gamma-ray bands. The resulting value p1 = 1.7−2, is
close to the canonical nonrelativistic value of 2 (see, e.g., Drury
1983). The location of the valley between the peaks is most sen-
sitive to the maximum γ factor of the electrons and the onset of
the IC peak. The maximum energies of the electrons are tuned
by the acceleration parameter ξ (see Eq. (1)), with an additional
constraint from the VHE gamma-ray spectrum. The depth of the
valley is additionally modified by the SSC component. By tuning
the compactness of the blob (i.e., varying the radius of the blob
within a factor of two while simultaneously keeping the same
total power in electrons that fixes the synchrotron and EC emis-
sion at roughly the same level), which affects the SSC emission,
we fit the depth of the dip in periods B, C, and D. As in all the
studied energy bands the emission during the whole investigated
period was significantly larger than the historical measurements,
we neglect the possible low-state emission of the source in the
20 Also called light crossing time scale.
A163, page 12 of 19
MAGIC Collaboration: MWL campaign on QSO B1420+326
Table 5. Parameters of core A0 and knot K20.
MJD Knot S R Θ a P EVPA
(1) (2) (3) (4) (5) (6) (7) (8)
58916 A0 0.90± 0.06 0.0 – 0.029± 0.007 3.5± 0.3 110± 6
58979 A0 0.52± 0.06 0.0 – 0.023± 0.007 1.5± 0.5 −20± 8
59006 A0 0.44± 0.05 0.0 – 0.066± 0.010 0.8± 0.3 78± 11
58916 K20 0.46± 0.04 0.073± 0.018 118± 9 0.088± 0.014 6.5± 0.4 113± 7
58979 K20 0.26± 0.03 0.118± 0.024 120± 8 0.096± 0.017 5.7± 0.3 114± 6
59006 K20 0.10± 0.02 0.151± 0.015 117± 8 0.067± 0.015 8.3± 0.5 91± 6
Notes. Columns of table are: (1) epoch of observation in MJD; (2) designation of knot; (3) flux density, S , in Jy; (4) distance from the core, R,
in mas; (5) position angle with respect to the core, Θ, in degrees; (6) FWHM size of knot, a, in mas; (7) degree of polarization, P, in %; and (8)
position angle of polarization, EVPA, in degrees. For model parameters, 1σ uncertainties are provided.
Fig. 11. Separation of K20 from the core with time (red triangles)
according to modeling; the red line represents a linear approximation
of the motion. Blue circles and the dashed line correspond to the light
curve of K20 at 43 GHz; the black vertical line indicates the time of the
VHE event.
modeling. The broadband spectra obtained from the modeling
are compared with measured ones in Fig. 12. The corresponding
evolution of EED is shown in Fig. 13. Parameters of the model-
ing are shown in Table 6.
It should be stressed that, mainly due to fixing the Doppler
factor, these latter parameters are not uniquely determined (see,
e.g., Ahnen et al. 2017a). For example, assuming δ and Γ moti-
vated by VLBI measurements, a similar fit to the data can be
achieved, albeit with the size of the emission region compressed
by about an order of magnitude in periods A, C, and D. However,
they allow us to trace the relative evolution of the parameters
with the assumption that the beaming did not change during the
flare. The gamma-ray spectrum in the modeling is reproduced
by almost the EC process alone. The high spectral variability of
the X-ray emission is naturally explained by three processes that
contribute to it: mainly the EC and SSC emission, with a par-
tial component from the synchrotron radiation from the highest
energies during period C.
Because of its simplicity, the modeling has some caveats.
The radio emission is underestimated because of pronounced
synchrotron self absorption in the model curve. Such emission
is often attributed to a larger scale jet, rather than the blob. How-
ever, it should be noted that during the 2020 high state, the radio
emission was at a higher level than in historical measurements,
and also the emission showed some variability, and therefore it
should be also at least partially associated with the high state.
It is plausible that, because of the evolution of the flare, low-
energy electrons from the blob escape to the large-scale jet with-
out being cooled completely. Such an enhancement of the EED
in the large-scale jet during the high state would explain higher
radio emission. In addition, such escape of high-energy particles
along the jet would naturally explain the new component appear-
ing in the VLBI follow-up of the flare.
In period C (and partially also in period D) the X-ray data
show a clear valley between the two peaks. The shape of the
low-energy (i.e., falling) part of the valley in this period is not
fully reproduced by the model. This part of the SED is strongly
dependent on the shape of the high-energy tail of the EED which
is also constrained by the highest energy gamma-rays. The sim-
plifications in the modeling (homogeneous region, no nonsta-
tionary processes modifying EED within the considered period,
resulting sharp cut-off of the EED) do not allow us to realisti-
cally model the full shape of the valley. The SED in this period
might also be affected by fast variability.
The slope of the spectrum is not accurately modeled in all the
cases. In particular, in period D (and partially also in period C)
the optical range would require softer electron distribution than
the gamma-ray range. In period C, the NIR data are slightly over-
shooting the model suggesting a softer spectrum. This might be
connected with the fast variability of the source, or with addi-
tional emission of the population of partially cooled particles
from the previous phases of the flare.
Despite those caveats, it is interesting to see that the obtained
model parameters provide a self-consistent description of the
main features of the emission during different phases of the flare.
Comparing period B to period A, modeling suggests a compres-
sion of the emission region coincident with the increase of the
minimum Lorentz factor of the electrons and a mild increase
of the magnetic field density and the total energy stored in
electrons.
The increased γmin factor needed in the modeling (see also
Katarzyn´ski et al. 2006) might point to a two-stage accelera-
tion process. First, injection of particles with such a minimum
Lorentz factor has to occur, for example because of accelera-
tion in a small potential drop due to reconnection of magnetic
fields (see, e.g., Lazarian et al. 2015; Comisso & Sironi 2019).
A second process (e.g., Fermi second-order acceleration) would
then boost the particles to a power-law spectrum up to maxi-
mal gamma factors of γmax. The acceleration coefficient ξ has
a rather small value, of the order of 10−7. The value of ξ in
the modeling increases by an order of magnitude in period C
in order to explain the VHE gamma-ray emission. Such small
values are needed to saturate the acceleration process with EC
energy losses in order not to overly increase γmax and in turn not
A163, page 13 of 19
A&A 647, A163 (2021)
108 1011 1014 1017 1020 1023 1026
[Hz]
10 14
10 13
10 12
10 11
10 10
10 9
10 8
F
[e
rg
cm
2
s
1 ]
Synch.
SSC
EC DT
Total
SSDC
A
A U.L.
108 1011 1014 1017 1020 1023 1026
[Hz]
10 14
10 13
10 12
10 11
10 10
10 9
10 8
F
[e
rg
cm
2
s
1 ]
Synch.
SSC
EC DT
Total
SSDC
B
B U.L.
108 1011 1014 1017 1020 1023 1026
[Hz]
10 14
10 13
10 12
10 11
10 10
10 9
10 8
F
[e
rg
cm
2
s
1 ]
Synch.
SSC
EC DT
Total
SSDC
C
C U.L.
108 1011 1014 1017 1020 1023 1026
[Hz]
10 14
10 13
10 12
10 11
10 10
10 9
10 8
F
[e
rg
cm
2
s
1 ]
Synch.
SSC
EC DT
Total
SSDC
D
D U.L.
Fig. 12. Multiwavelength SED of QSO B1420+326 in the four periods: A – before the flare, B – optical flare, C – VHE gamma-ray flare, D –
after the flare and archival data (gray). Different radiation processes are shown with different line styles: dotted lines – synchrotron, dashed – SSC,
dot-dashed – EC, solid – sum of components. Model lines are corrected for EBL absorption according to Domínguez et al. (2011).
100 101 102 103 104
102
103
2 d
n e
/d
[e
l.
cm
3 ]
A
B
C
D
100 101 102 103 104
1052
1053
1054
2 d
N
e/d
[e
l.]
A
B
C
D
Fig. 13. Evolution of the EED in the frame of the blob: energy density (left) and total energy integrated over the blob (right). Different line colors
represent different stages of the flare.
to overshoot the X-ray emission by synchrotron component. As
shown in the modeling, such acceleration would still be efficient
enough to explain the observed optical and HE flare. A natu-
ral explanation for such low values of ξ is a second-order Fermi
process with nonrelativistic scattering centers accelerating elec-
trons in the emission region. During period D, the VHE emission
requires a further small increase of the ξ parameter coincident
with lower magnetic field in order not to overshoot the soft X-ray
flux with the synchrotron component.
5. Discussion and conclusions
Observations of QSO B1420+326 with MAGIC during the
enhanced state allowed us to add a new member to the sparse
family of FSRQs emitting in the VHE range. The observations
were performed during an impressive flare, with the flux of both
SED peaks enhanced by about two orders of magnitude with
respect to the low state. Monitoring observations of the source
and a massive MWL campaign provided us a dataset that was
used to trace the evolution of the EED during the flare. Inter-
A163, page 14 of 19
MAGIC Collaboration: MWL campaign on QSO B1420+326
Table 6. Parameters used for the modeling.
Period δ rb [cm] ξ B [G] U′e [1048 erg] p1 γmin p2 γbreak γmax u′e [erg cm−3] u′e/u′B
A 40 6.16 × 1016 0.3 × 10−7 0.70 1.18 1.7 1 2.7 63 6900 1.2 × 10−3 0.06
B 40 3.70 × 1016 0.3 × 10−7 0.95 1.76 1.8 15 2.8 104 8000 8.3 × 10−3 0.23
C 40 3.08 × 1016 3.0 × 10−7 0.83 2.12 2.0 10 3.0 125 23700 17.3 × 10−3 0.63
D 40 3.08 × 1016 6.0 × 10−7 0.55 2.35 2.0 10 3.0 125 27300 19.2 × 10−3 1.6
Notes. Columns are: Doppler factor δ (Γ = δ is assumed), co-moving size of the emission region rb, acceleration parameter ξ, magnetic field B,
total energy of electron U′e, EED: slope before the break: p1, minimum Lorentz factor γmin, slope after the break p2, the Lorentz factor of the break
γbreak, maximum Lorentz factor γmax, electron energy density u′e, energy equipartition u
′
e/u
′
B. Free parameters of the model and derived parameters
are put on the left and right side of the double vertical line, respectively.
estingly, the synchrotron spectrum in the optical range during
the flare (in particular period B) is hard. Comparing to histori-
cal data, both low- and high-energy peaks shifted by at least two
orders of magnitude in frequency, such a large shift being rare for
a FSRQ object. Shifts of the peaks towards higher energies dur-
ing high states is a behavior commonly observed in a sister class
of objects, the BL Lacs. The spectra are Compton dominated,
which is typical for FSRQs. However, the dominance during the
peak of the flare is just a factor of a few.
Similarly to other VHE-detected FSRQs, we get a satisfac-
tory description of the two broad peaks of the SED as the syn-
chrotron and EC on the DT radiation field. The valley between
the peaks is well constrained by the X-ray data, and we use it to
track the changes of the compactness of the blob during the evo-
lution of the high state. The modeling scenario is self-consistent
in the sense that the shape of the EED is determined by the bal-
ance of acceleration and cooling processes. The variability of
the emission between different phases of the enhanced state is
explained mainly by a combination of variations of the com-
pactness of the emission region, the minimal injection energy
of electrons, and the increase in the acceleration parameters. In
addition, to achieve a satisfactory fit, coincident small variations
of the magnetic field, total energy stored in electrons, and an
injection slope have been assumed.
The optical spectroscopy observations revealed a prominent
MgII line that does not show flux variability exceeding the uncer-
tainties of the measurements. We explain this as the line being
produced within a canonical BLR, which means that it has a
much longer timescale of variability. Therefore, the MgII line
is a good proxy for estimating the accretion disk luminosity
at 2 × 1046 erg s−1. Additionally, a broad FeII bump has been
observed, with the luminosity increasing with the increase of
the optical continuum emission. The fast variability of the FeII
bump suggests that it originates in a much smaller region (possi-
bly located close to the jet axis) than the regular BLR. Moreover,
as the flux of FeII bump correlates with the synchrotron contin-
uum, the bump should be produced farther from the black hole
in the vicinity of the jet. Additionally, the shift of the bump to
the blue side could be explained if it is produced in a wind sur-
rounding the jet. This suggests a possible interaction of the jet
with a FeII-emitting cloud.
Optical polarimetry that started a few days after optical peak
shows a very low polarization and smooth EVPA rotation. This
makes QSO B1420+326 another FSRQ in which VHE gamma-
ray emission is detected contemporaneous to EVPA rotations.
Intriguingly, the VLBI observations in the follow-up observa-
tions of the flare show an emission of a superluminal radio knot
contemporaneous with the high gamma-ray state. A similar asso-
ciation of VHE gamma-ray emission, with EVPA rotation, and
VLBI component ejection has previously been suggested for
another FSRQ, PKS 1510−089 (Ahnen et al. 2017b).
The detection of QSO B1420+326 in the VHE gamma-ray
range and the extensive monitoring campaign during this event
add another piece to the puzzle of the origin of the highest
energy emission of FSRQ objects. It is interesting that some
of the emission features associated with these observations of
QSO B1420+326 are similar to observations of other FSRQs, in
particular of PKS 1510−089, so far the most thoroughly studied
FSRQ in the VHE range. QSO B1420+326 might be a cousin of
PKS 1510−089, twice as distant but intrinsically more luminous.
Acknowledgements. We would like to dedicate this paper to the memory of
our friend and colleague, Dr. Valeri Larionov (1950−2020), who enthusias-
tically contributed to this and many other projects aimed at understanding
blazars. We would like to thank the Instituto de Astrofísica de Canarias for
the excellent working conditions at the Observatorio del Roque de los Mucha-
chos in La Palma. The financial support of the German BMBF and MPG;
the Italian INFN and INAF; the Swiss National Fund SNF; the ERDF under
the Spanish MINECO (FPA2017-87859-P, FPA2017-85668-P, FPA2017-82729-
C6-2-R, FPA2017-82729-C6-6-R, FPA2017-82729-C6-5-R, AYA2015-71042-
P, AYA2016-76012-C3-1-P, ESP2017-87055-C2-2-P, FPA2017-90566-REDC);
the Indian Department of Atomic Energy; the Japanese ICRR, the University
of Tokyo, JSPS, and MEXT; the Bulgarian Ministry of Education and Science,
National RI Roadmap Project DO1-268/16.12.2019 and the Academy of Finland
grant nr. 320045 is gratefully acknowledged. This work was also supported by
the Spanish Centro de Excelencia “Severo Ochoa” SEV-2016-0588 and SEV-
2015-0548, the Unidad de Excelencia “María de Maeztu” MDM-2014-0369 and
the “la Caixa” Foundation (fellowship LCF/BQ/PI18/11630012), by the Croa-
tian Science Foundation (HrZZ) Project IP-2016-06-9782 and the University
of Rijeka Project 13.12.1.3.02, by the DFG Collaborative Research Centers
SFB823/C4 and SFB876/C3, the Polish National Research Centre grant UMO-
2016/22/M/ST9/00382 and by the Brazilian MCTIC, CNPq and FAPERJ. The
Fermi-LAT Collaboration acknowledges generous ongoing support from a num-
ber of agencies and institutes that have supported both the development and
the operation of the LAT, as well as scientific data analysis. These include the
National Aeronautics and Space Administration and the Department of Energy
in the United States; the Commissariat à l’Energie Atomique and the Centre
National de la Recherche Scientifique/Institut National de Physique Nucléaire et
de Physique des Particules in France; the Agenzia Spaziale Italiana and the Isti-
tuto Nazionale di Fisica Nucleare in Italy; the Ministry of Education, Culture,
Sports, Science and Technology (MEXT), High Energy Accelerator Research
Organization (KEK), and Japan Aerospace Exploration Agency (JAXA) in
Japan; and the K. A. Wallenberg Foundation, the Swedish Research Council, and
the Swedish National Space Board in Sweden. Additional support for science
analysis during the operations phase is gratefully acknowledged from the Isti-
tuto Nazionale di Astrofisica in Italy and the Centre National d’Etudes Spatiales
in France. This work was performed in part under DOE Contract DE-AC02-
76SF00515. This publication makes use of data obtained at the Metsähovi Radio
Observatory, operated by Aalto University in Finland. This research has made
use of data from the OVRO 40 m monitoring program Richards et al. (2011)
which is supported in part by NASA grants NNX08AW31G, NNX11A043G, and
NNX14AQ89G and NSF grants AST-0808050 and AST-1109911. This study
was based in part on observations conducted using the 1.8 m Perkins Telescope
Observatory (PTO) in Arizona, which is owned and operated by Boston Univer-
sity. The research at Boston University was supported in part by NASA Fermi
GI program grants 80NSSC17K0649, 80NSSC19K1504, and 80NSSC19K1505.
A163, page 15 of 19
A&A 647, A163 (2021)
We thank the ASAS-SN team for making their data publicly available. The
VLBA is an instrument of the National Radio Astronomy Observatory. The
National Radio Astronomy Observatory is a facility of the National Science
Foundation operated by Associated Universities, Inc. This work made use of the
Lowell Discovery Telescope (formerly Discovery Channel Telescope) at Lowell
Observatory. Lowell is a private, nonprofit institution dedicated to astrophys-
ical research and public appreciation of astronomy and operates the LDT in
partnership with Boston University, the University of Maryland, the University
of Toledo, Northern Arizona University, and Yale University. We acknowledge
support by Bulgarian National Science Fund under grant DN18-10/2017 and
National RI Roadmap Projects DO1-277/16.12.2019 and DO1-268/16.12.2019
of the Ministry of Education and Science of the Republic of Bulgaria.
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1 Instituto de Astrofísica de Canarias and Dpto. de Astrofísica,
Universidad de La Laguna, 38200 La Laguna, Tenerife, Spain
2 Università di Udine and INFN Trieste, 33100 Udine, Italy
3 National Institute for Astrophysics (INAF), 00136 Rome, Italy
4 ETH Zürich, 8093 Zürich, Switzerland
5 Institut de Física d’Altes Energies (IFAE), The Barcelona Institute
of Science and Technology (BIST), 08193 Bellaterra, Barcelona,
Spain
A163, page 16 of 19
MAGIC Collaboration: MWL campaign on QSO B1420+326
6 Japanese MAGIC Group: Institute for Cosmic Ray Research
(ICRR), The University of Tokyo, Kashiwa 277-8582, Chiba, Japan
7 Technische Universität Dortmund, 44221 Dortmund, Germany
8 Croatian MAGIC Group: University of Zagreb, Faculty of Electrical
Engineering and Computing (FER), 10000 Zagreb, Croatia
9 IPARCOS Institute and EMFTEL Department, Universidad Com-
plutense de Madrid, 28040 Madrid, Spain
10 Centro Brasileiro de Pesquisas Físicas (CBPF), 22290-180 URCA,
Rio de Janeiro, RJ, Brazil
11 University of Lodz, Faculty of Physics and Applied Informatics,
Department of Astrophysics, 90-236 Lodz, Poland
12 Università di Siena and INFN Pisa, 53100 Siena, Italy
13 Deutsches Elektronen-Synchrotron (DESY), 15738 Zeuthen, Ger-
many
14 Università di Padova and INFN, 35131 Padova, Italy
15 INFN MAGIC Group: INFN Sezione di Torino and Università degli
Studi di Torino, 10125 Torino, Italy
16 Max-Planck-Institut für Physik, 80805 München, Germany
17 Università di Pisa and INFN Pisa, 56126 Pisa, Italy
18 Universitat de Barcelona, ICCUB, IEEC-UB, 08028 Barcelona,
Spain
19 Armenian MAGIC Group: A. Alikhanyan National Science Labora-
tory, Yerevan, Armenia
20 Centro de Investigaciones Energéticas, Medioambientales y Tec-
nológicas, 28040 Madrid, Spain
21 INFN MAGIC Group: INFN Sezione di Bari and Dipartimento
Interateneo di Fisica dell’Università e del Politecnico di Bari, 70125
Bari, Italy
22 Croatian MAGIC Group: University of Rijeka, Department of
Physics, 51000 Rijeka, Croatia
23 Universität Würzburg, 97074 Würzburg, Germany
24 Finnish MAGIC Group: Finnish Centre for Astronomy with ESO,
University of Turku, 20014 Turku, Finland
25 Departament de Física, and CERES-IEEC, Universitat Autònoma de
Barcelona, 08193 Bellaterra, Spain
26 Armenian MAGIC Group: ICRANet-Armenia at NAS RA, Yerevan,
Armenia
27 Croatian MAGIC Group: University of Split, Faculty of Electri-
cal Engineering, Mechanical Engineering and Naval Architecture
(FESB), 21000 Split, Croatia
28 Croatian MAGIC Group: Josip Juraj Strossmayer University of Osi-
jek, Department of Physics, 31000 Osijek, Croatia
29 Japanese MAGIC Group: Department of Physics, Kyoto University,
606-8502 Kyoto, Japan
30 Japanese MAGIC Group: Department of Physics, Tokai University,
Hiratsuka 259-1292, Kanagawa, Japan
31 Saha Institute of Nuclear Physics, HBNI, 1/AF Bidhannagar, Salt
Lake, Sector-1, Kolkata 700064, India
32 Inst. for Nucl. Research and Nucl. Energy, Bulgarian Academy of
Sciences, 1784 Sofia, Bulgaria
33 Finnish MAGIC Group: Astronomy Research Unit, University of
Oulu, 90014 Oulu, Finland
34 Croatian MAGIC Group: Rud¯er Boškovic´ Institute, 10000 Zagreb,
Croatia
35 INFN MAGIC Group: INFN Sezione di Perugia, 06123 Perugia,
Italy
36 INFN MAGIC Group: INFN Roma Tor Vergata, 00133 Roma, Italy
37 Now at University of Innsbruck, Innsbruck, Austria
38 Also at Port d’Informació Científica (PIC), 08193 Bellaterra,
Barcelona, Spain
39 Also at Dipartimento di Fisica, Università di Trieste, 34127 Trieste,
Italy
40 Max-Planck-Institut für Physik, 80805 München, Germany
41 Also at INAF Trieste and Dept. of Physics and Astronomy, Univer-
sity of Bologna, Bologna, Italy
42 Japanese MAGIC Group: Institute for Cosmic Ray Research
(ICRR), The University of Tokyo, Kashiwa 277-8582, Chiba, Japan
43 ASI Space Science Data Center, Via del Politecnico, snc., 00133
Rome, Italy
44 INFN – Roma Tor Vergata, Via della Ricerca Scientifica, 1., 00133
Rome, Italy
45 INAF – IRA Bologna, Via P. Gobetti 101, 40129 Bologna, Italy
46 Istituto Nazionale di Fisica Nucleare, Sezione di Perugia, 06123
Perugia, Italy
47 Naval Research Laboratory, Washington, DC 20375, USA
48 INAF – Istituto di Astrofisica e Planetologia Spaziali, Via Fosso del
Cavaliere, 100, 00133 Rome, Italy
49 Physical Research Laboratory, Ahmedabad, India
50 Institute of Astronomy and NAO, Bulgarian Academy of Sciences,
72 Tsarigradsko Shose Blvd., 1784 Sofia, Bulgaria
51 Department of Astronomy, Faculty of Physics, University of Sofia,
1164 Sofia, Bulgaria
52 Università di Siena, 53100 Siena, Italy
53 Instituto Nacional de Astrofisica, Optica & Electronics Tonantzintla,
Puebla, Mexico
54 Aalto University Metsähovi Radio Observatory, Metsähovintie 114,
02540 Kylmälä, Finland
55 Aalto University Department of Electronics and Nanoengineering,
PO Box 15500, 00076 Aalto, Finland
56 Finnish Centre for Astronomy with ESO (FINCA), University of
Turku, Vesilinnantie 5, 20014 Turku, Finland
57 Institute for Astrophysical Research, Boston University, 725 Com-
monwealth Avenue, Boston, MA 02215, USA
58 Astronomical Institute, St. Petersburg University, Universitetskij Pr.
28, Petrodvorets 198504, St. Petersburg, Russia
59 Pulkovo Observatory, St. Petersburg, Russia
60 Crimean Astrophysical Observatory, Crimea
61 Owens Valley Radio Observatory, California Institute of Technol-
ogy, Pasadena, CA 91125, USA
62 Institute of Astrophysics, Foundation for Research and Technology-
Hellas, 71110 Heraklion, Greece
63 Department of Physics, Univ. of Crete, 70013 Heraklion, Greece
64 Departamento de Astronomía, Universidad de Chile, Camino El
Observatorio 1515, Las Condes, Santiago, Chile
65 CePIA, Departamento de Astronomía, Universidad de Concepción,
Concepción, Chile
A163, page 17 of 19
A&A 647, A163 (2021)
Appendix A: X-ray short-term variations, spectral
fits and long-term multiwavelength behaviour
A.1. XMM-Newton and Swift-XRT short-term variations
We investigated short-term variability of the source in UV and
X-rays by analyzing Swift and XMM-Newton data. Correcting for
instrumental artifacts and pile-up, no significant (>3-σ) increase
of count rate was observed in the background-subtracted consec-
utive XRT observation segments.
We also searched for significant (>3-σ) changes of magni-
tude between two consecutive UVOT exposures collected with
the same filter at the same epoch. This resulted in two events
observed in w1 filter on 2020 January 19 (MJD = 58867) and
31 (MJD = 58879), with a change of 0.16 mag in 27.6 ks and
0.41 mag in 34.4 ks, respectively.
Moreover, we produced a light curve of the XMM-Newton pn
count rate with a bin of 500 s (Fig. A.1). The light curve shows
only moderate variability, with the count rate varying between
2.08 and 2.86 cps. The fractional variability (see Vaughan et al.
2003 for details) is 0.064± 0.005. A larger variability was
observed in the second part of the observation, in particular an
increase of ∼15% of the count rate at the time of the highest
peak.
A.2. Swift-XRT spectral fits
In Table A.1 we report the results of analysis of all Swift-XRT
observations, including both historical observations and the ones
performed during the high state of the source shown in detail in
Fig. 1.
A.3. Long-term behaviour
In Fig. A.2 we report the monitoring observations of
QSO B1420+326 in order to put the flaring state of 2020 January
and February in the context of the long-term behavior of the
source. It is clear that the flaring period has been unique in the
Fermi-LAT dataset of this source. Similarly, the radio flux during
the flare is also unique compared to previous measurements. On
the other hand, there is a prior report of a similar magnitude opti-
cal flare on 2016 March 11 (MJD = 57458.5) in the ASAS-SN
monitoring data (Stanek et al. 2017), albeit without a HE coun-
terpart. Past X-ray data are unfortunately too sparse (and biased
by Target of Opportunity observations) to make conclusions as
to the typical behavior of the source.
Fig. A.1. XMM-Newton EPIC pn light curve over 0.3−10 keV with 500 s
bins of QSO B1420+326.
Table A.1. Log and fitting results of Swift-XRT observations of QSO B2 1420+326 using a PL model.
Date MJD Net exposure time Photon index Flux0.3−10 keV
(UT) (s) (ΓX) (10−12 erg cm−2 s−1)
2018-02-22 58171.915197 2872 1.49± 0.22 1.99± 0.37
2018-12-12 58464.132329 1818 1.51± 0.25 2.76± 0.54
2018-12-14 58466.253218 2110 1.37± 0.18 4.15± 0.65
2018-12-16 58468.386537 1943 1.50± 0.21 3.22± 0.54
2019-06-25 58659.880024 1975 1.89± 0.21 3.18± 0.45
2019-06-27/29 58663 2008 1.25± 0.24 3.07± 0.64
2019-07-12 58676.328225 1593 2.33± 0.15 7.11± 0.69
2019-07-17 58681.609499 1983 2.17± 0.15 4.83± 0.48
2019-07-22 58686.722614 1885 1.45± 0.22 2.98± 0.55
2019-12-13 58830.935420 2023 1.44± 0.13 7.85± 0.87
2020-01-02 58850.318109 2495 1.33± 0.17 5.47± 0.75
2020-01-05 58853.943061 1666 1.38± 0.27 4.09± 0.87
2020-01-08 58856.665569 1880 1.69± 0.20 3.73± 0.55
2020-01-11 58859.830600 1983 1.81± 0.17 5.54± 0.70
2020-01-19 58867.533771 1626 1.87± 0.15 5.88± 0.65
2020-01-21 58870.011577 1546 1.94± 0.16 6.00± 0.65
2020-01-24 58872.175468 1783 1.84± 0.14 6.57± 0.65
2020-01-25 58873.171868 1798 2.10± 0.14 7.12± 0.69
2020-01-27 58875.267506 1806 1.73± 0.19 4.42± 0.61
2020-01-28 58876.193806 2035 1.47± 0.18 4.80± 0.67
2020-01-31 58879.685875 1231 1.68± 0.19 5.88± 0.83
2020-02-01 58880.114444 1681 1.58± 0.16 5.87± 0.77
2020-02-05 58884.201846 2025 1.49± 0.15 6.30± 0.75
2020-02-10 58889.473095 2417 1.57± 0.18 3.50± 0.48
Notes. Fluxes are corrected for the Galactic absorption.
A163, page 18 of 19
MAGIC Collaboration: MWL campaign on QSO B1420+326
Fermi-LAT F >0.1GeV
55000 55500 56000 56500 57000 57500 58000 58500 59000
0
5
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15
]
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s
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cm
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[10
Fermi-LAT index
55000 55500 56000 56500 57000 57500 58000 58500 59000
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Ph
ot
on
 in
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x
Swift XRT 0.3-10 keV
55000 55500 56000 56500 57000 57500 58000 58500 59000
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]
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er
g 
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F[
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55000 55500 56000 56500 57000 57500 58000 58500 590001
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W2-band
55000 55500 56000 56500 57000 57500 58000 58500 59000
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[m
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V-band
55000 55500 56000 56500 57000 57500 58000 58500 59000
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ASAS-SN V&g
ASAS-SN V&g U.L.
UVOT V
OVRO 
55000 55500 56000 56500 57000 57500 58000 58500 59000
Time [MJD]
0
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F 
[Jy
]
Fig. A.2. Long-term MWL light curve of QSO B1420+326 (see titles and legends of individual panels). The gray-shaded region shows the flaring
period in the beginning of 2020.
A163, page 19 of 19