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A&A 594, A25 (2016)
DOI: 10.1051/0004-6361/201526803
c
ESO 2016
Astronomy
&
Astrophysics
Planck 2015 results Special feature
Planck 2015 results
XXV. Diffuse low-frequency Galactic foregrounds
Planck Collaboration: P. A. R. Ade95 , N. Aghanim64, M. I. R. Alves105,11,64 , M. Arnaud79, M. Ashdown75,6, J. Aumont64, C. Baccigalupi93 ,
A. J. Banday105,11, R. B. Barreiro70 , J. G. Bartlett1,72, N. Bartolo33,71 , E. Battaner107,108, K. Benabed65,104, A. Benoît62 , A. Benoit-Lévy27,65,104,
J.-P. Bernard105,11 , M. Bersanelli36,52, P. Bielewicz89,11,93, J. J. Bock72,13, A. Bonaldi73 , L. Bonavera70, J. R. Bond10, J. Borrill16,99 ,
F. R. Bouchet65,97, F. Boulanger64 , M. Bucher1, C. Burigana51,34,53, R. C. Butler51 , E. Calabrese101, J.-F. Cardoso80,1,65, A. Catalano81,78,
A. Challinor67,75,14, A. Chamballu79,18,64 , R.-R. Chary61, H. C. Chiang30,7, P. R. Christensen90,39, S. Colombi65,104, L. P. L. Colombo26,72,
C. Combet81, F. Couchot77, A. Coulais78 , B. P. Crill72,13 , A. Curto70,6,75, F. Cuttaia51, L. Danese93, R. D. Davies73, R. J. Davis73 , P. de Bernardis35 ,
A. de Rosa51, G. de Zotti48,93 , J. Delabrouille1, J.-M. Delouis65,104, F.-X. Désert58, C. Dickinson73, ∗, J. M. Diego70, H. Dole64,63 , S. Donzelli52,
O. Doré72,13, M. Douspis64 , A. Ducout65,60, X. Dupac41 , G. Efstathiou67, F. Elsner27,65,104, T. A. Enßlin85, H. K. Eriksen68, E. Falgarone78,
J. Fergusson14, F. Finelli51,53, O. Forni105,11, M. Frailis50 , A. A. Fraisse30, E. Franceschi51 , A. Frejsel90, S. Galeotta50, S. Galli74 , K. Ganga1,
T. Ghosh64, M. Giard105,11, Y. Giraud-Héraud1, E. Gjerløw68, J. González-Nuevo22,70, K. M. Górski72,110 , S. Gratton75,67, A. Gregorio37,50,57 ,
A. Gruppuso51, J. E. Gudmundsson102,91,30 , F. K. Hansen68, D. Hanson87,72,10 , D. L. Harrison67,75, G. Helou13 , S. Henrot-Versillé77 ,
C. Hernández-Monteagudo15,85, D. Herranz70 , S. R. Hildebrandt72,13, E. Hivon65,104, M. Hobson6, W. A. Holmes72, A. Hornstrup19, W. Hovest85,
K. M. Huffenberger28, G. Hurier64 , A. H. Jaffe60, T. R. Jaffe105,11, W. C. Jones30, M. Juvela29 , E. Keihänen29, R. Keskitalo16, T. S. Kisner83,
R. Kneissl40,8, J. Knoche85, M. Kunz20,64,3, H. Kurki-Suonio29,46, G. Lagache5,64 , A. Lähteenmäki2,46, J.-M. Lamarre78 , A. Lasenby6,75,
M. Lattanzi34,54, C. R. Lawrence72 , J. P. Leahy73, ?, R. Leonardi9, J. Lesgourgues66,103, F. Levrier78, M. Liguori33,71, P. B. Lilje68,
M. Linden-Vørnle19, M. López-Caniego41,70 , P. M. Lubin31 , J. F. Macías-Pérez81 , G. Maggio50, D. Maino36,52 , N. Mandolesi51,34,
A. Mangilli64,77, M. Maris50 , D. J. Marshall79, P. G. Martin10, E. Martínez-González70, S. Masi35 , S. Matarrese33,71,43, P. McGehee61,
P. R. Meinhold31 , A. Melchiorri35,55, L. Mendes41 , A. Mennella36,52, M. Migliaccio67,75 , S. Mitra59,72, M.-A. Miville-Deschênes64,10,
A. Moneti65, L. Montier105,11 , G. Morgante51, D. Mortlock60, A. Moss96 , D. Munshi95, J. A. Murphy88 , F. Nati30 , P. Natoli34,4,54 ,
C. B. Netterfield23, H. U. Nørgaard-Nielsen19 , F. Noviello73 , D. Novikov84, I. Novikov90,84, E. Orlando109, C. A. Oxborrow19 , F. Paci93 ,
L. Pagano35,55, F. Pajot64, R. Paladini61, D. Paoletti51,53, B. Partridge45 , F. Pasian50, G. Patanchon1, T. J. Pearson13,61, M. Peel73, O. Perdereau77 ,
L. Perotto81, F. Perrotta93, V. Pettorino44, F. Piacentini35, M. Piat1, E. Pierpaoli26 , D. Pietrobon72, S. Plaszczynski77, E. Pointecouteau105,11 ,
G. Polenta4,49, G. W. Pratt79, G. Prézeau13,72, S. Prunet65,104 , J.-L. Puget64, J. P. Rachen24,85, W. T. Reach106, R. Rebolo69,17,21, M. Reinecke85 ,
M. Remazeilles73,64,1, C. Renault81, A. Renzi38,56 , I. Ristorcelli105,11, G. Rocha72,13 , C. Rosset1, M. Rossetti36,52, G. Roudier1,78,72,
J. A. Rubiño-Martín69,21, B. Rusholme61 , M. Sandri51, D. Santos81 , M. Savelainen29,46, G. Savini92, D. Scott25, M. D. Seiffert72,13 ,
E. P. S. Shellard14 , L. D. Spencer95, V. Stolyarov6,100,76, R. Stompor1, A. W. Strong86, R. Sudiwala95, R. Sunyaev85,98 , D. Sutton67,75,
A.-S. Suur-Uski29,46, J.-F. Sygnet65, J. A. Tauber42, L. Terenzi94,51, L. Toffolatti22,70,51 , M. Tomasi36,52, M. Tristram77, M. Tucci20, J. Tuovinen12,
G. Umana47, L. Valenziano51, J. Valiviita29,46 , F. Van Tent82 , M. Vidal73, P. Vielva70, F. Villa51, L. A. Wade72, B. D. Wandelt65,104,32, R. Watson73,
I. K. Wehus72,68, A. Wilkinson73, D. Yvon18, A. Zacchei50, and A. Zonca31
(Affiliations can be found after the references)
Received 22 June 2015 /Accepted 20 April 2016
ABSTRACT
We discuss the Galactic foreground emission between 20 and 100 GHz based on observations by Planck and WMAP. The total intensity in this
part of the spectrum is dominated by free-free and spinning dust emission, whereas the polarized intensity is dominated by synchrotron emission.
The Commander component-separation tool has been used to separate the various astrophysical processes in total intensity. Comparison with radio
recombination line templates verifies the recovery of the free-free emission along the Galactic plane. Comparison of the high-latitude Hαemission
with our free-free map shows residuals that correlate with dust optical depth, consistent with a fraction (≈30%) of Hαhaving been scattered by
high-latitude dust. We highlight a number of diffuse spinning dust morphological features at high latitude. There is substantial spatial variation in
the spinning dust spectrum, with the emission peak (in Iν) ranging from below 20 GHz to more than 50 GHz. There is a strong tendency for the
spinning dust component near many prominent H ii regions to have a higher peak frequency, suggesting that this increase in peak frequency is
associated with dust in the photo-dissociation regions around the nebulae. The emissivity of spinning dust in these diffuse regions is of the same
order as previous detections in the literature. Over the entire sky, the Commander solution finds more anomalous microwave emission (AME) than
the WMAP component maps, at the expense of synchrotron and free-free emission. This can be explained by the difficulty in separating multiple
broadband components with a limited number of frequency maps. Future surveys, particularly at 5–20 GHz, will greatly improve the separation by
constraining the synchrotron spectrum. We combine Planck and WMAP data to make the highest signal-to-noise ratio maps yet of the intensity of
the all-sky polarized synchrotron emission at frequencies above a few GHz. Most of the high-latitude polarized emission is associated with distinct
large-scale loops and spurs, and we re-discuss their structure. We argue that nearly all the emission at 40◦>l>−90◦is part of the Loop I structure,
and show that the emission extends much further in to the southern Galactic hemisphere than previously recognised, giving Loop I an ovoid rather
than circular outline. However, it does not continue as far as the “Fermi bubble/microwave haze”, making it less probable that these are part of
the same structure. We identify a number of new faint features in the polarized sky, including a dearth of polarized synchrotron emission directly
correlated with a narrow, roughly 20◦long filament seen in Hαat high Galactic latitude. Finally, we look for evidence of polarized AME, however
many AME regions are significantly contaminated by polarized synchrotron emission, and we find a 2σupper limit of 1.6% in the Perseus region.
Key words. diffuse radiation – ISM: general – radiation mechanisms: general – radio continuum: ISM – polarization – local insterstellar matter
∗Corresponding authors: C. Dickinson, clive.dickinson@manchester.ac.uk; J. P. Leahy, j.p.leahy@manchester.ac.uk
Article published by EDP Sciences A25, page 1 of 45
A&A 594, A25 (2016)
1. Introduction
Diffuse Galactic radio emission consists of a number of distinct
components that emit via different emission mechanisms (no-
tably synchrotron, free-free, and spinning dust). There is con-
siderable interest in understanding these components in order
to subtract foregrounds cleanly from cosmic microwave back-
ground (CMB) data and as a probe of the physics of the interstel-
lar medium (ISM) and Galactic structure. The separation of the
diffuse foregrounds into their separate constituent components is
extremely difficult, since their spectra (except for free-free emis-
sion) are not well-known, and they have comparable intensities
at microwave frequencies (Leach et al. 2008).
At low frequencies (.10 GHz), synchrotron radiation from
electrons spiralling in the Galactic magnetic field dominates the
sky. At frequencies around 1 GHz the spectral index (T∝νβ)
is β≈ −2.7, while at higher frequencies it appears to steepen
to β≈−3.0. This broadly agrees with theoretical calculations,
such as those from the GALPROP1code (Orlando & Strong
2013); however, a detailed comparison has not been made up
to now, due to difficulties in the component separation process.
Synchrotron radiation is intrinsically highly polarized and is a
strong polarized CMB foreground up to around 100 GHz. Free-
free emission emits over a range of radio frequencies and can
be significant up to about 100 GHz owing to its flatter spectral
index (β=−2.1). Anomalous microwave emission (AME) is
an additional foreground that has been detected at frequencies
10–60 GHz (Leitch et al. 1997;de Oliveira-Costa et al. 2004;
Finkbeiner 2004;Davies et al. 2006;Planck Collaboration XX
2011;Planck Collaboration XXI 2011;Planck Collaboration
Int. XV 2014) above the expected levels of synchrotron and
free-free emission. The most plausible origin for AME is elec-
tric dipole radiation from tiny, rapidly spinning dust grains
(Draine & Lazarian 1998), which provides an excellent fit to the
data, particularly from well-studied molecular clouds (Planck
Collaboration XX 2011). However, the diffuse emission at high
latitudes has still to be definitively identified, and other possibil-
ities exist, such as magneto-dipole radiation from thermal fluctu-
ations of magnetized grains (Draine & Lazarian 1999;Liu et al.
2014). The polarization of AME is of great interest to CMB cos-
mologists, since a significant level of polarization could hamper
measurements of Bmodes. Current measurements place upper
limits on the polarization of AME at a few per cent or below
(Dickinson et al. 2011;Macellari et al. 2011;Rubiño-Martín
et al. 2012a).
In the 2013 Planck2results based on spectral fitting, the com-
ponent separation procedure did not achieve a separation of the
distinct low-frequency components, instead lumping these into
a single low-frequency foreground (Planck Collaboration XII
2014). However, other methods have been used to do this un-
der specific assumptions; Planck Collaboration Int. XII (2013)
used correlated component analysis in the southern Gould Belt
region where synchrotron emission is relatively smooth and
weak, while Planck Collaboration Int. XXII (2015) used tem-
plate fitting to separate components that were correlated with
1http://galprop.stanford.edu;
http://sourceforge.net/projects/galprop/
2Planck (http://www.esa.int/Planck) is a project of the
European Space Agency (ESA) with instruments provided by two sci-
entific consortia funded by ESA member states and led by Principal
Investigators from France and Italy, telescope reflectors provided
through a collaboration between ESA and a scientific consortium led
and funded by Denmark, and additional contributions from NASA
(USA).
each spatial template. In the 2015 analysis (Planck Collaboration
X 2016) we include the WMAP (Bennett et al. 2013a) and
408 MHz (Haslam et al. 1982) all-sky maps to enable a separa-
tion of synchrotron, free-free, and AME. The separation is based
on fitting a spectral model to each sky pixel using Commander, a
Gibbs sampling code (Eriksen et al. 2008).
The current Planck data release also includes polarization
information. At angular scales larger than 1◦the polarized emis-
sion is dominated by synchrotron and thermal dust radiation.
Synchrotron radiation is more important at low frequencies and
dominates the polarized Planck maps at 30 and 44 GHz. This
synchrotron emission comes mainly from the Galactic plane and
also from filamentary structures that can extend over 100◦across
the sky. While these “loops” and “spurs” have long been known
in total intensity (Quigley & Haslam 1965;Berkhuijsen et al.
1971;Berkhuijsen 1971), their pattern is clarified by the new
polarization maps. Vidal et al. (2015) present an analysis of the
filaments using WMAP data. They have catalogued them, stud-
ied the polarization angle distribution and polarized spectral in-
dices, and presented a model to explain the origin of some of
them. Here, with the new Planck polarization data, we expand on
this, identifying new polarized features and understanding better
the already known ones, thanks to the improved signal-to-noise
ratio (S/N) of the data.
We begin in Sect. 2by summarizing the data sets employed,
including a brief discussion of some of the most important sys-
tematic effects. We then take a first look at the intensity fore-
grounds in Sect. 3by employing a constrained internal linear
combination (ILC) algorithm to remove the CMB and free-free
emission components whose spectra are well known. In Sect. 4
we then discuss each of the low frequency foreground compo-
nents (AME, free-free, and synchrotron emission) as determined
by the Commander algorithm. We use these new maps to investi-
gate the distribution and spectra of these foregrounds. In Sect. 5
we use WMAP/Planck polarization maps to study the distribu-
tion of the low frequency polarized foregrounds. We make a high
S/N synchrotron map by combining the WMAP/Planck data into
a single product, assuming a power-law model for synchrotron
emission. We discuss in detail the large-scale features in the
low frequency polarized sky including the well-known loops and
spurs. Other results will be published in separate articles, includ-
ing the diffuse synchrotron power spectrum and further analysis
of AME polarization constraints from Galactic clouds. We con-
clude in Sect. 6.
2. Data
Table 1lists the primary radio data sets used in our analysis, with
properties that are important here. We note that brightness tem-
peratures are in the Rayleigh-Jeans convention unless otherwise
stated (WMAP and Planck data are natively in CMB thermody-
namic temperature). We now briefly discuss these data sets and
particular issues that are relevant to our analysis.
2.1. WMAP and Planck data
The primary data used in the analysis are the 9-year WMAP
maps (Bennett et al. 2013a) and the 2015 release of Planck
maps (Planck Collaboration I 2016;Planck Collaboration
II 2016;Planck Collaboration VIII 2016) from the Low
Frequency Instrument (LFI) and High Frequency Instrument
(HFI). Details of the map preparation, including smooth-
ing to 1◦full-width half-maximum (FWHM) resolution and
correction of residual offsets and gain errors, are given in
A25, page 2 of 45
Planck Collaboration: Planck 2015 results. XXV.
Table 1. Summary of radio data sets used, with selected properties.
Map Instrument νa
0η∆T(ν0)bσIcw0dw1dw2dw3d
[GHz] [µK]
Haslam Effelsberg/Jodrell/Parkes 0.408 1.000 3 ×106. . . . . . . . . . . .
Kband WMAP 22.8 0.987 6.00 . . . 1.007 . . . 1.549
30 GHz Planck LFI 28.4 0.979 2.49 1.013 . . . 3.585 −0.485
Ka band WMAP 33.0 0.972 4.44 ... ... ... −0.788
Qband WMAP 40.6 0.958 3.86 ... ... ... −1.342
44 GHz Planck LFI 44.1 0.951 2.80 . . . . . . −9.571 −2.457
Vband WMAP 60.8 0.910 4.32 . . . . . . . . . . . .
70 GHz Planck LFI 70.4 0.881 2.22 . . . . . . . . . . . .
Wband WMAP 93.5 0.802 5.14 . . . . . . . . . . . .
100 GHz Planck HFI 100 0.777 0.86 . . . . . . . . . . . .
143 GHz Planck HFI 143 0.592 0.37 −1.052 −1.088 6.194 3.648
217 GHz Planck HFI 217 0.334 0.51 . . . . . . . . . . . .
353 GHz Planck HFI 353 0.075 1.69 0.039 0.038 −0.208 −0.126
545 GHz Planck HFI 545 0.006 0.51e. . . . . . . . . . . .
857 GHz Planck HFI 857 6 ×10−50.48e. . . . . . . . . . . .
Notes. (a)Reference frequency. (b)Conversion factor from differential thermodynamic temperature to Rayleigh-Jeans brightness temperature.
(c)Median rms per beam for the 1◦smoothed Stokes Imaps. (d)Weights for the linear combination images shown in Figs. 1a and 2a (w0), Fig. 2b
(w1), Fig. 3a (w2), and Fig. 3b (w3). (e)Units MJy sr−1.
Planck Collaboration X (2016). The products are at HEALPix
Nside =256. In addition, for polarization analysis we have pre-
pared maps at 2◦resolution, using the same methodology as for
the 1◦maps. Component separation depends critically on the ac-
curacy of the input data, and so we briefly review data quality
issues here.
Thermal noise. This is accurately characterized in the Planck
and WMAP data. For both missions, the noise is quite vari-
able over the sky. We have calculated the thermal noise in our
smoothed maps via propagation of errors, assuming the noise is
independent between pixels at our full resolution. After smooth-
ing, the noise in nearby pixels is strongly correlated.
Confusion. Confusion noise denotes the uncertainties caused
by faint compact sources within each beam. It has been esti-
mated from source counts made at much higher angular reso-
lution. To some extent, confusion is automatically fitted out as
part of component separation. This would be exactly true if the
source had a spectrum that could be well fitted by our model, and
was non-variable. In fact, typical sources responsible for confu-
sion in our maps are blazars, characterized by strong variability
and “flat” (β≈ −2) spectra, which are similar to free-free emis-
sion, but may contain multiple peaks (e.g., Planck Collaboration
XIII 2011). Hence a fraction of the confusion is an effective
noise component in our fitting.
Calibration. Uncertainties for both WMAP and Planck due
to calibration are about 0.2% for most frequencies and this
contributes negligibly to our error budget. The two highest-
frequency HFI channels have raw uncertainties of .10% (Planck
Collaboration VIII 2016), and corrections were derived as part
of the component separation fit; these uncertainties have very
little impact on the low-frequency components discussed in this
paper. The global calibration uncertainty of the Haslam et al.
(1982) map is about 10%, but this only impacts our estimate
for the break in the synchrotron spectrum. There is also an ad-
ditional source of error that arises from the broad bandwidths
of the Planck and WMAP detectors. We can take these into ac-
count by applying “colour corrections” that depend on the spec-
tral response of the instrument and the spectrum of the sky sig-
nal (Jarosik et al. 2011;Planck Collaboration II 2014;Planck
Collaboration IX 2014;Planck Collaboration II 2016;Planck
Collaboration VII 2016). For Planck LFI these corrections are
typically 1–3% while they are 5–10% for HFI. The residual er-
rors from such corrections are likely to be .1%, especially for
the LFI channels of interest here. We implement WMAP colour
corrections using the same approach as for the LFI, allowing for
the secular drift of the centre frequency (Bennett et al. 2013a)
and the small bandpass shifts derived in Planck Collaboration
X(2016). Uncertainties due to residual beam asymmetry in the
smoothed WMAP/Planck data will be, relative to other uncer-
tainties, minimal because we have used the deconvolved WMAP
data and the additional smoothing to 1◦FWHM resolution will
symmeterize the beam.
Offsets. Neither Planck nor WMAP measured the absolute zero
level in total intensity, and so there is an arbitrary offset in each
map. There is also a small but obvious residual dipole in dif-
ferences between Planck and WMAP maps. In the analysis in
Sect. 3we use the approach of Wehus et al. (2016) that uses
T–Tplots to estimate the dipole and monopole amplitudes for
each map (exact values were updated based on the 2015 release
Planck maps); the approach used for the Commander analysis is
described in Planck Collaboration X (2016). The zero levels in
the polarization maps (Stokes Qand U) are much better defined,
and are recovered with negligible error by Planck. In contrast,
two large angular scale modes in the WMAP Qand Umaps
are sensitive to zero-level uncertainties in the raw data. These
“poorly constrained modes” were down-weighted by the WMAP
team in their power spectrum analysis (Bennett et al. 2013a), but
this approach is not available when attempting to model the spec-
trum pixel-by-pixel. They are strong enough to seriously affect
the absolute polarized brightness over most of the sky, except for
the regions of strongest polarized intensity, such as the Galactic
plane and the North Polar Spur.
A25, page 3 of 45
A&A 594, A25 (2016)
Polarization leakage. Various processes cause leakage of to-
tal intensity into the polarization signal, notably mismatch be-
tween the beamshapes for the two polarizations in each horn,
and bandpass mismatch (e.g., Leahy et al. 2010). The WMAP
scanning strategy allows a rather accurate correction for these ef-
fects, but a more indirect approach is needed for Planck (Planck
Collaboration II 2016) and the corrections applied are believed
to be accurate to a few tenths of a percent of the total intensity.
This is a significant problem along the Galactic plane, where the
fractional polarization is very low and so the residual bandpass
leakage error is relatively high. Beam mismatch leakage is also
strong on the plane, due to the steep intensity gradients there. But
for most of the sky, the residual leakage is essentially negligible.
Unmodelled emission. This includes the Sunyaev-Zeldovich
(SZ) effect over most of the sky3, most spectral lines, and the
spectrum of extragalactic point sources. Hundreds of transitions
are known in the part of the spectrum of interest here, but over
most of the sky they are expected to be very weak compared to
the continuum in our broad bands. However, they are not always
negligible: rich line spectra are generated in photo-dissociation
regions (PDRs) surrounding H ii regions, and also in the X-ray
dominated region around the Galactic centre (Takekawa et al.
2014). While this spectral contamination is certainly unusual,
there are no blind spectral line surveys in most of our bands that
would allow us to set reliable upper limits to line contamina-
tion. The Commander fit actually includes line emission (CO and
HCN) in four of our bands.
2.2. Ancillary data
Below 20 GHz, several radio maps are available. Reich et al.
(2004) have presented a preliminary map at 1.4 GHz, but the
combination of instrumental artefacts and calibration uncertain-
ties in the survey do not meet our stringent requirements on the
error budget, which also caused us to drop a number of Planck
data sets as described in Planck Collaboration X (2016). We do
use the 408-MHz map of Haslam et al. (1982), as re-processed
by Remazeilles et al. (2015) to reduce instrumental striping and
provide better point-source subtraction than previous versions.
The noise in the Haslam et al. (1982) map is not characterized in
detail, but in any case is much less than residual scanning arte-
facts. We assume an rms error of 1 K, including both effects. We
also assume a 10% uncertainty in the overall intensity scale.
The Galactic synchrotron spectrum shows significant curva-
ture below 10 GHz, while free-free absorption sets in for bright
Hii regions below a few hundred MHz. We therefore do not con-
sider maps at frequencies below 408 MHz.
We use the full-sky dust-corrected Hαmap of Dickinson
et al. (2003) to compare with our free-free solution at high lat-
itudes where dust absorption is relatively small. The fraction of
dust lying between us and the Hα-emitting gas is assumed to be
1/3 ( fd=0.33). We also use the Finkbeiner (2003) Hαmap to
test the results. At low latitudes, we make use of the radio recom-
bination line (RRL) survey recovered from the HIPASS data at
1.4 GHz (Alves et al. 2015). This survey detects RRLs from ion-
ized gas along the narrow Galactic plane with a resolution of 150.
Although of limited sensitivity, it provides the radio equivalent
of Hαmaps without absorption along the line of sight.
3The largest SZ signals, due to the Coma and Virgo clusters, are ex-
plicitly fitted by Commander, but this is only feasible because other fore-
ground components are very faint in these regions.
We use the IRIS reprocessed far infrared (FIR) map from
the IRAS survey at 100, 60, 25, and 12 µm (Miville-Deschênes
& Lagache 2005). The intensity scale is good to 13% and in-
strumental noise is effectively insignificant after smoothing to
1◦resolution. We also use the H imap from the LAB survey
(Kalberla et al. 2005).
We also make use of the ROSAT diffuse sky survey maps.
The original images provided by the Max-Planck-Institut für
Extraterrestrische Physik (MPE)4are stored as six panels that
between them cover the whole sky, with slight overlaps between
adjacent panels. The pixel size is 120. The 0.1–2 keV energy
range is divided into six bands. The publicly available FITS im-
ages are derived from the analysis by Snowden et al. (1997) of
the 1990–1991 ROSAT survey observations, and do not include
the make-up data collected in 1997 and presented by Freyberg &
Egger (1999); consequently they show narrow strips of missing
data. After correcting erroneous field centres in the MPE FITS
headers, we found the pixel coordinates in the original images of
HEALPix (Górski et al. 2005) pixels at Nside =2048, so that the
original pixels are highly oversampled. Nearest-neighbour inter-
polation was used for placing the intensity and uncertainty val-
ues into the HEALPix grid. We averaged the data and uncertain-
ties onto an Nside =256 pixel grid (pixel size 13.
07, comparable
to the original), weighting each raw pixel by the area contributed
to the output HEALPix pixel. We then merged the HEALPix maps
of the six original panels into a single all-sky map, using simple
averaging for pixels covered by two or more panels5.
3. A first look at the low-frequency foregrounds
using constrained internal linear combinations
Of all the emission mechanisms contributing to the cm-
wavelength sky, the CMB is the most accurately measured,
since the data are calibrated directly on the CMB dipole, and
it is also the only component with a precisely known spec-
trum. It can therefore be straightforwardly eliminated by con-
structing linear combinations of the individual frequency maps.
Therefore, as a first look, we apply a linear combination to the
WMAP/Planck maps to investigate the range of spectral indices
and morphologies at frequencies around 20–50 GHz. Our ap-
proach (Appendix A) is similar to, but simpler and more limited
than, standard ILC techniques.
Figure 1a shows the total foreground emission in the Planck
30 GHz band (reference frequency of 28.4GHz), constructed via
Tforeground =w30T30 +w143 T143 +w353T353 (1)
where the subscripts refer to channel frequencies, the weights
wiare given in Table 1and, as always in this section, are nor-
malized to return a β=−3 power law unchanged. We use the
Planck HFI 143 GHz band as our primary model of the CMB,
since it has the lowest noise and is dominated by only two com-
ponents, the CMB itself and thermal dust emission. The latter
is cleaned to adequate precision for display purposes using the
HFI 353 GHz map. The dust is modelled as a uniform mod-
ified blackbody spectrum with βd=1.51 and Td=19.6 K
4http://www.xray.mpe.mpg.de/rosat/survey/sxrb/12/ass.
html. Although these web pages display GIF images including the
1997 make-up observations, the downloadable FITS images include
only the 1990–91 data.
5All-sky HEALPix maps of the six ROSAT bands created as de-
scribed here are available at http://www.jb.man.ac.uk/research/
cosmos/rosat/
A25, page 4 of 45
Planck Collaboration: Planck 2015 results. XXV.
(a)
(b)
(c) (d)
Fig. 1. a) CMB-nulled map at 28.4 GHz, constructed using the ILC method (see text), where in this map, β=−3 emission is unchanged by
construction; b) 408 MHz map (Remazeilles et al. 2015), strongly dominated by synchrotron emission; c) 545 GHz Planck map, strongly dominated
by thermal dust emission; and d) Hαmap (Dickinson et al. 2003). An asinh(I) colour scale is used for all images, where Iis the intensity in the
units indicated. The coordinate along the colour bar is linear in Inear zero and becomes logarithmic (ln(2I)) at high intensity. All maps are at 1◦
resolution.
(Planck Collaboration Int. XXII 2015). This combination frac-
tionally diminishes the amplitude of free-free emission, but only
by 0.7%, which is negligible for display purposes.
Figure 1also shows our ancillary maps at 408 MHz (trac-
ing synchrotron emission), Hα(tracing free-free emission),
and the Planck 545 GHz map (tracing thermal dust emission).
Features of all these templates appear in the 28.4 GHz fore-
ground map: the brightest features along the Galactic plane are
mostly free-free emission; at high latitudes in the northern hemi-
sphere the North Polar synchrotron spur is prominent; while
obvious dust-related features include the ρOph complex, the
Chameleon/Musca complex of cold clouds arcing around the
South Celestial Pole, the Polaris Flare, and the comet-shaped
R CrA molecular cloud. In addition to the diffuse components,
there are thousands of point sources. In fact, at this resolution
the map is strongly confusion-limited (Fig. 2); the median ther-
mal noise is 2.5 µK, but the rms in patches a few degrees across
is about 10 µK, even in regions far from the plane where cat-
alogued sources have been avoided. Consequently a signal is
detected significantly in almost every pixel, and only a few SZ
decrements (notably from the Coma Cluster) give a negative map
intensity.
After the CMB, the next best determined spectrum is free-
free emission. Assuming local thermodynamic equilibrium, the
spectral shape depends, very weakly, on a single parameter, the
electron temperature Te; we use the approximation of Draine
(2011) detailed in Sect. 4.1 below. For plausible values, 3000 <
Te<15 000 K, the ratio of observed brightness between our 30
and 44 GHz bands varies only between 2.52 and 2.49 (includ-
ing the effect of colour corrections). We can therefore null this
component, taking Te=7500 K, by differencing CMB-corrected
maps at 28.4 and 44.1 GHz. Figure 3a shows the result.
The drawback of this procedure is that the thermal noise is
amplified by a factor of 11, partly because the differencing in-
troduces noise from the 44 GHz image, and partly due to the
renormalization. The situation can be considerably improved by
including WMAP data: Fig. 3b shows a similar analysis that in-
cludes the three lowest-frequency WMAP bands, where we max-
imize the S/N of our reference β=−3 spectrum as well as
eliminating CMB, free-free, and thermal dust. In this case the
S/N is 5.5 times worse than in the original 28.4 GHz map; the
better performance is partly due to the larger frequency ratio
(22.8/44.1 vs. 28.4/44.1), which reduces the unwanted cancel-
lation of the synchrotron emission.
Emission with a power-law spectrum steeper than free-free,
i.e., β < −2.1, will show up as positive in the free-free nulled
maps, and these show (albeit with higher noise) the character-
istic synchrotron and AME features noted in our description of
Fig. 1above. Most of the extragalactic sources are nearly nulled,
since their spectrum is similar to free-free; those remaining are
mostly optically-thin synchrotron sources, but sources that var-
ied between the WMAP and Planck mean epochs also show up
(as negative residuals if they were brighter during the Planck
observations).
In this double difference image the assumed thermal dust
spectrum begins to have a significant impact; if we use βd=1.66
and Td=19.0 K (Planck Collaboration Int. XIV 2014) the in-
ner Galactic plane becomes about 20% brighter. Given that Tdis
known to vary significantly between the inner and outer Galaxy
A25, page 5 of 45
A&A 594, A25 (2016)
Fig. 2. Left:a) south polar region of the CMB-nulled 28.4 GHz map. Overplotted circles are the positions of 30-GHz sources from the second
Planck Catalogue of Compact Sources (PCCS2, Planck Collaboration XXVI 2016). The 90% completeness limit for the catalogue (426 mJy)
corresponds to a peak brightness of ∆T≈50 µK at the 1◦resolution of this map. Right:b) same region in a CMB-nulled version of the WMAP
K-band map (22.8 GHz). The diffuse arc in the range 30◦<l<150◦is part of Loop II. The colour scales are linear here.
(a)
(b)
Fig. 3. Maps with CMB, free-free, and thermal dust emission nulled,
scaled so that a β=−3 power-law spectrum is unchanged in amplitude.
a) Planck-only map (28.4, 44.1, 143, 353 GHz), scaled to 28.4 GHz.
b) Combined Planck and WMAP. Weights used are given in Table 1. An
asinh colour scale is used and an offset of 80 µKCMB has been subtracted
to enhance the contrast.
(e.g., Planck Collaboration XI 2014), more sophisticated mod-
elling is required to obtain the most from the data.
Nevertheless, the ILC approach does reveal an important
complexity of the low-frequency foreground spectrum in a
straightforward way, namely as prominent extended regions of
negative intensity, notably coincident with the Sh2-27 H ii region
at (l,b)=(6◦,24◦), and in the southern part of the Gum nebula,
around (l,b)=(257◦,−14◦). These are regions with excess flux
compared to the free-free spectrum near 40 GHz. They can be
fitted with a spinning dust spectrum shifted to peak at higher
frequencies than standard models, but still within the range of
plausible interstellar parameters (e.g., Planck Collaboration XX
2011;Planck Collaboration Int. XV 2014).The Sh2-27 spinning
dust is likely to be associated with the famous translucent cloud
in which interstellar molecules were first discovered, as absorp-
tion lines in the spectrum of the ionizing star ζOph (Adams
1941). In fact, many bright H ii regions show negative residuals,
e.g., the California nebula (l,b)=(161◦,−13◦), Orion A and B
(l,b)=(210◦,−19◦) and (207◦,−17◦), and many of the bright
nebulae along the Galactic plane.
These residuals are not due to errors in our assumed Teval-
ues; pushing Teas high as 15 000 K makes almost no difference
to the spectrum and so leaves all these features in place. They are
also insensitive to changes in estimated colour corrections be-
tween Planck Collaboration II (2014) and Planck Collaboration
II (2016). They become less prominent for steeper assumed ther-
mal dust spectra (βd), but this is mainly due to the impact of βd
on the underlying Galactic ridge emission noted above.
The systematic negative residuals suggest that high-
frequency spinning dust may often be associated with the PDRs
around H ii regions. Dobler & Finkbeiner (2008b) also iden-
tified a potential high-frequency spinning dust component in
WMAP data that was correlated with Hαfrom the Warm Ionized
Medium. An alternative explanation in at least some cases (in
particular the Galactic centre) might be significant line emission
in the 39–46 GHz region common to WMAP Q and the LFI 44-
GHz bands; however well-known lines in this band, such as SiO
(1–0) and methanol masers, do not seem to be bright enough to
cause the effect (Jordan et al. 2015; Jordan, priv. comm.). The
A25, page 6 of 45
Planck Collaboration: Planck 2015 results. XXV.
implications of the wide range in AME peak frequency are dis-
cussed in Sect. 4.2.
From this preliminary analysis we can draw two further
lessons. First, separating components with similar spectra makes
great demands on sensitivity; even separating free-free emis-
sion (β=−2.1) from synchrotron (β≈ −3) has dramati-
cally reduced our effective S/N. The synchrotron spectral in-
dex varies by only a few tenths across the sky, and to map it
with useful precision we need a substantial gain in sensitivity.
This has been achieved mainly by working at low resolution
(Fuskeland et al. 2014;Vidal et al. 2015); in the current paper
we use a fixed template for the synchrotron spectrum (Sect. 4.3
below).
Second, low-frequency foregrounds have complicated spec-
tra: at a minimum, we must solve for the amplitudes of free-free
emission, synchrotron emission, and AME, and for the latter the
peak frequency and spectral width. However, even including the
WMAP data, we have only five frequency maps between 22 and
44 GHz to solve for these five parameters. With no spare degrees
of freedom, we would be susceptible to both errors in the as-
sumed spectral model and any artefacts in the data. We therefore
need to use maps at higher and lower frequencies.
Above 44 GHz we have the Planck maps at 70 and 100 GHz
and WMAP V- and W-bands (60 and 94 GHz). There is signif-
icant line contamination in the W- and 100 GHz-bands, which
has to be included in the component separation. In all these
maps the strongest foreground component is the low-frequency
tail of the thermal dust emission, for which it is necessary to fit
at least the temperature Tdand probably the emissivity index βd
as well. As a result, including these higher frequency maps gives
us more data points but also require us to solve for more spectral
parameters. Moreover, since the foreground emission is weak-
est at these frequencies, the maps have relatively poor S/N. The
Commander approach discussed in the next section is our best
attempt to date to handle these complexities.
4. Total intensity foregrounds with Commander
The analysis of Sect. 3demonstrated the complexity of the low-
frequency foreground spectrum, but for that very reason was un-
able to separate the various diffuse foreground components that
contribute to the maps. A more sophisticated method is required,
and numerous algorithms exist that utilize the spectral and/or
morphological information in the data. One of the most powerful
algorithms is parametric fitting via Gibbs sampling, which has
been implemented in the Commander code (Eriksen et al. 2008).
The details of the specific implementation to these data are de-
scribed in Planck Collaboration X (2016). In this section, we
discuss the models for the low-frequency components used by
Commander, and compare the derived foreground products with
previous results such as WMAP component-separation products.
We also compare the results with expectations, either based on
theory or extrapolations from ancillary data.
4.1. Free-free
The free-free radio continuum brightness can be estimated by the
use of recombination lines. The most widely used is the Hαline
(λ656.28 nm), which traces warm ionized gas and is proportional
to the emission measure (EM), in the same way as free-free radio
continuum (Dickinson et al. 2003). Thus Hαis a reliable estima-
tor provided that the gas is in local thermodynamic equilibrium
(LTE), that there is no significant dust absorption along the line-
of-sight, that there is no significant scattered component of Hα
light, and that the electron temperature is known. The relation-
ship between the free-free brightness Tff
b[µK] and Hαintensity
IHα[R] for the optically thin limit is given by
Tff
b
IHα
=1512 T0.517
4100.029/T4ν−2.0
GHz gff,(2)
where T4is the electron temperature in units of 104K, νGHz the
frequency in GHz, and gffis the Gaunt factor, which takes into
account quantum mechanical effects. In the pre-factor we have
included a factor of 1.08 to account for the He ii contribution
that adds to the free-free continuum6.Draine (2011) provides an
approximation to gffthat is accurate to within about 1% in the
Planck frequency range:
gff(ν, Te)=ln
exp
5.960 −√3
πln νGHzT−3/2
4
+e
.(3)
At high Galactic latitudes, the absorption by dust is a small effect
(.0.1 mag), and can be corrected for to first-order. The electron
temperature is known to vary throughout the Galaxy, but has typ-
ical values of (7500 ±1000)K in the local diffuse ISM (Shaver
et al. 1983;Paladini et al. 2004;Alves et al. 2012). This corre-
sponds to a theoretical radio-to-Hαratio of (11.1±0.9) µKR−1
at 22.8 GHz, assuming Te=(7500 ±1000) K. However, previ-
ous template fitting results have indicated that the measured ra-
tio is lower than this value, at ≈7–10 µKR−1(Davies et al. 2006;
Ghosh et al. 2012).
This apparent discrepancy can be resolved with a lower elec-
tron temperature (Te≈4000 K), although this is outside the
range expected for our position in the Galaxy. Dust absorption
cannot explain it since it has the opposite effect: it would in-
crease the derived Hαemissivity resulting in an even lower ratio.
The most popular explanation is that a fraction of the observed
Hαintensity is from light scattered by dust grains (assuming the
gas is close to being in LTE). Previous estimates of this frac-
tion suggested relatively low values (around 10%), while more
recent analyses indicate that as much as half of the high latitude
Hαintensity could be scattered (Witt et al. 2010). Other anal-
yses give values around the 10–20% level (Wood & Reynolds
1999;Brandt & Draine 2012;Barnes et al. 2015). We attempt to
estimate the scattered fraction in Sect. 4.1.1. We begin by com-
paring our derived free-free amplitude with the observed Hα.
We note that there are essentially no constraints on Tefrom
the Commander solution, which is based on variations in the
free-free spectral index, except in the inner Galactic plane (see
Sect. 4.1.2).
Figure 4(top panel) shows a T–Tplot of the free-free am-
plitude against the Hαintensity for the Gum nebula region
(which is known to be dominated by free-free emission), as de-
fined in Fig. 5. A very good correlation can be seen, which in-
dicates the robustness of the Commander free-free amplitude.
The best-fitting slope is (8.9±0.9) µKR−1. This corresponds
to an electron temperature of (5200 ±900) K, which is consis-
tent for the temperature measured around the same region using
Hydrogen RRL by Woermann et al. (2000). We also measure
the free-free-Hαcorrelation using a full-sky map, masking the
more dusty regions (Fig. 4, bottom panel). The mask was de-
fined by ignoring all pixels with an optical depth at 353 GHz,
τ353, larger than 1.5×10−5(this corresponds to an extinction at
6We note that in Planck Collaboration X (2016), the equation for the
free-free continuum brightness does not take into account this additional
contribution from helium.
A25, page 7 of 45
A&A 594, A25 (2016)
Fig. 4. Top:T–Tplot of the free-free amplitude at 22.8 GHz against the
Hαintensity [R] in the Gum Nebula (see Fig. 5for the location of the
region). Bottom:T–Tplot for 85% of the sky. Regions that show large
dust absorption have been masked out, mostly within |b|.7◦.
the Hαwavelength of A(Hα)=0.54 mag, assuming a reddening
value RV=3.1; Planck Collaboration XI 2014), which results
in a masked area of approximately 15% of the sky, mostly on
the Galactic plane. In the region studied, we find a best-fitting
slope between the Commander free-free map and the Hαmap of
(8.0±0.8) µKR−1. This is close to the value for the Gum nebula,
and is consistent with previous analyses of the free-free-to-Hα
ratio at this frequency (Banday et al. 2003;Davies et al. 2006;
Ghosh et al. 2012). This gives us additional confidence that the
component separation has worked relatively well for the free-
free component. Assuming the theory is correct, this lower value
for the free-free to Hαratio suggests a lower value for Te, of
around 4500 K. However, it is likely that a fraction of the Hαin-
tensity is from scattering of Hαlight from the Galactic ridge by
dust grains.
4.1.1. High latitude Hαscattering
In this section we estimate the fraction, fscatt, of Hαlight that is
scattered by dust grains. This is believed to be responsible for
Fig. 5. Commander free-free full-sky map at 22.8 GHz. The box shows
the region we selected in the Gum nebula for the scatter plot shown in
Fig. 4.
the low free-free to Hαratio discussed in the previous section,
and could reduce the intrinsic Hαintensity by up to ≈50% (Witt
et al. 2010). To do this, we must constrain Te, since there is a
one-to-one degeneracy between Teand fscatt. For high latitude
sight-lines away from the Galactic plane, we can assume that
the bulk of the gas is at the local value Te=(7500 ±1000) K,
thus breaking the degeneracy. We initially ignore variations in
dust absorption, since we have masked out the regions that are
most affected. We test the effect of this assumption by making a
nominal correction for dust absorption.
First, we use the theoretical proportionality constant that re-
lates the Hαintensity with the Commander free-free emission,
(11.1±0.9) µKR−1, to scale the free-free template to Rayleigh
units. We note that the uncertainty accounts for a ±1000 K un-
certainty in Te. We subtract this template from the Hαmap to
obtain the Hαemission that does not come from in-situ recom-
bination, i.e., originating from Hαlight scattered by dust grains.
The top panel in Fig. 6shows the residual map, with the pix-
els masked out shown in grey and the colour scale being trun-
cated to ±2R, which saturates the brightest regions closer to
the Galactic plane and the Gould Belt. A feature of this residual
map is the blue central region, which resembles the microwave
haze noted by Finkbeiner (2004), Dobler & Finkbeiner (2008a)
and Planck Collaboration Int. IX (2013). This is not unexpected,
since the haze was discovered by assuming that Hαis an accurate
free-free template. This is an indication that component separa-
tion is ambiguous in this complex region, since mixed emission
from the Gould Belt system makes the separation between AME,
free-free, and synchrotron particularly difficult. There may also
be component separation issues in other regions of the sky at
a lower level, which could have a significant impact on these
residuals.
However, we can still set an upper limit for the amount of
scattered Hαlight, based on the value of the free-free to Hα
ratio measured over the sky. The mean value at 22.8 GHz that
we measure over most of the sky is (8.0±0.8) µKR−1. From
this, if we assume that the difference relative to the theoretical
value of (11.1±0.9) µKR−1is due to an excess of Hαemis-
sion from scattering, we can estimate the scattering fraction to
be fscatt =(28 ±12)%. This is an average value for the high lat-
itude sky, so individual regions with different electron tempera-
ture might be present a level higher than this. It is consistent with
previous estimates of around 20% of scattered light on average
at high latitudes (Wood & Reynolds 1999;Witt et al. 2010;Dong
& Draine 2011;Brandt & Draine 2012;Barnes et al. 2015). If we
repeat the calculation using a dust-corrected Hαmap, assuming
A25, page 8 of 45
Planck Collaboration: Planck 2015 results. XXV.
Fig. 6. Top: residual Hαemission after subtraction of the scaled free-
free component. Bottom: scatter plot of the residual Hαintensity against
the thermal dust optical depth at 353 GHz outside the mask shown
above, corresponding to A(Hα).0.53 (τ353 <1×10−5). The best-
fitting linear relation is shown as a dashed line.
that 1/3 of the dust lies in front of the Hα-emitting gas (see
Dickinson et al. 2003), we find fscatt =(36 ±12)%.
Our constraints are clearly not very tight, and are dependent
on our assumption on the mean electron temperature and on the
level of dust absorption. Given the relatively large uncertainties,
an electron temperature of about 5000 K would be enough to
bring the ratios within the 1σuncertainty of the observed range.
We now explore the correlation between the residual Hα
emission and the dust optical depth at 353 GHz. Scattering by
dust grains should produce a positive correlation between these
two maps. We plot this in the bottom panel of Fig. 6, where for
low values of the opacity, τ353 .1×10−5, the correlation is in-
deed positive. The large dispersion of the points in the figure is a
consequence of the noise and variations of electron temperature
at high Galactic latitude. In spite of that, a significant correlation
is observed. This suggests that the effect of scattered Hαlight
might be real and that its effects must be taken into account to
determine precise values of Te. But we remind the reader that
this could also partly be explained by the regions corresponding
to ≈0.5–1 R(at high Galactic latitudes, |b|&40◦, where the pos-
itive correlation is observed) having a lower than average value
for Te.
We do not find strong evidence for high latitude regions with
a scattered light fraction of around 50%. We also do not see an
anti-correlation of the free-free-to-Hαratio with the thermal dust
optical depth (a proxy for column density of dust), as would be
expected if there were a significant fraction of scattered light. We
hypothesize that the electron temperature at our exact Galactic
position could be slightly lower than expected (say 6000 K),
which would reduce fscatt to ≈10%.
4.1.2. Free-free emission in the plane
Along the Galactic plane, the Hαline is highly absorbed by dust
and thus fails to provide a reliable measure of the free-free emis-
sion (Dickinson et al. 2003). In this region of the sky, we can
use RRLs to estimate the thermal continuum emission and com-
pare with the Commander free-free solution. The integrated RRL
brightness is proportional to the free-free continuum in exactly
the same way as for the optical Hαrecombination line.
We use the data from the H iParkes All-Sky Survey
(HIPASS; Barnes et al. 2001) of the southern sky, which have
been re-analysed to extract RRLs at 1.4 GHz to produce the
first continguous RRL survey of the Galactic plane (Alves et al.
2015). The full survey covers the longitude range 196◦→0◦→
52◦and |b| ≤ 5◦, with an angular resolution of 14.
04. The inten-
sity calibration is good to better than 10%.
Figure 7shows maps of the free-free emission from the
Commander fit and from the RRLs. The free-free amplitude
is converted from the RRL temperature assuming a simple
model for Te, based on an assumed temperature gradient of
(500 ±100) K kpc−1with distance from the Galactic centre
(Alves et al. 2015). The Galactocentric distance at each longi-
tude is estimated based on the Fich et al. (1989) rotation curve
and using the central velocity of the RRLs. Obviously this is
a crude approximation, since the RRL is a line-of-sight inte-
gral weighted by n2
eand individual lines of sight could have a
range of values. We estimate that the electron temperatures for
a given line-of-sight are uncertain at the ±1000 K level, which
corresponds to an uncertainty in the free-free amplitude of up
to 20%. Thus the overall uncertainty on the free-free amplitude
from the RRL data is 25% at most. Nevertheless, there is an over-
all good level of agreement between the two estimates, both in
terms of morphology and amplitude. A combination of bright
compact sources (H ii regions) and diffuse emission within about
±1◦of the plane contribute, particularly in the inner Galaxy
(l=300◦→0◦→60◦). In the outer Galaxy (l<300◦) the
free-free emission is much weaker and only a few moderately
bright H ii regions can be seen.
The scatter plots of Fig. 8indicate that the two data sets
are highly correlated, which suggests that the component sep-
aration of the free-free emission from the other components has
been relatively successful. However, there are clearly changes in
slope from region to region. In general, the slopes are greater
than 1, indicating that the Commander free-free map is, on aver-
age, brighter that the predicted RRL map. The best-fitting slopes
are typically 50% higher than the predicted values.
Figure 9shows longitudinal slices across the Galactic plane
(b=0◦) at 1◦resolution. The bright H ii regions are clearly
visible as bright peaks. Virtually all the bright peaks in the
Commander free-free solution have a counterpart in the RRL
prediction. These peaks sit on diffuse emission that should be
subtracted before a comparison of peak temperatures is made;
large-scale (&4◦) emission will not be reliable in the RRL data
(Alves et al. 2015). Many of the peaks have comparable ampli-
tudes, while some of them are discrepant by up to a factor of
2. We note that the Galactic centre (l=0◦) is low in the RRL
map due to receiver saturation artefacts. The ratio of free-free
to RRL brightness for the brightest 6 peaks in the first region
A25, page 9 of 45
A&A 594, A25 (2016)
50 40 30 20 10 0
l [degree]
-4
-2
0
2
4
b [degree]
-1.0
7.0
15.0
50 40 30 20 10 0
l [degree]
-4
-2
0
2
4
b [degree]
-1.0
7.0
15.0
350 340 330 320 310
l
[degree]
-4
-2
0
2
4
b [degree]
-1.0
7.0
15.0
350 340 330 320 310
l [degree]
-4
-2
0
2
4
b [degree]
-1.0
7.0
15.0
300 290 280 270 260
l
[degree]
-4
-2
0
2
4
b [degree]
-1.0
7.0
15.0
300 290 280 270 260
l [degree]
-4
-2
0
2
4
b [degree]
-1.0
7.0
15.0
240 230 220 210 200
l [degree]
-4
-2
0
2
4
b [degree]
-1.0
4.5
10.0
240 230 220 210 200
l [degree]
-4
-2
0
2
4
b [degree]
-1.0
4.5
10.0
Fig. 7. Maps of the free-free emission along the Galactic plane, sepa-
rated into four different longitude ranges. For each l-range, the RRL and
Commander estimates are shown in the top and bottom panels, respec-
tively. The maps are in units of K at 1.4 GHz and at 1◦resolution. There
is a single contour for each pair of panels, set at a value corresponding to
the minimum temperature of the 2% brightest pixels of the RRL map,
corresponding to 7.4, 9.6, 4.4, and 0.4 K, from top to bottom, respec-
tively. The white pixels at the location of some bright H ii regions in the
maps of the first, third and fifth rows, correspond to saturated pixels in
the RRL survey.
(l=358◦→0◦→52◦, excluding the l=0◦peak) is 1.14±0.04,
while the next 10 brightest peaks have a ratio of 1.36 ±0.08 (the
uncertainty is the rms scatter of the ratios).
The discrepancy between the two free-free estimates on
bright H ii regions, which has been noted before (Alves et al.
2010;Planck Collaboration Int. XIV 2014;Planck Collaboration
Int. XXIII 2015), is presumably due to a combination of
effects: (i) differences in the electron temperature used to esti-
mate the free-free brightness; (ii) beam effects in the RRL map.
The RRL data are calibrated in the full beam scale, such that
the flux density of compact sources is underestimated by 10–
30% depending on their size, although the longitude profiles
of Fig. 9suggest that the discrepancy between the Commander
and RRL maps across individual sources does not depend on
their extent; and (iii) residual emission in the Commander map,
e.g., at (l,b)=(46.
◦8,−0.
◦3) associated with the HC30 supernova
remnant.
Compared to the brightest sources in each region, the
weaker diffuse emission appears to have an even steeper slope,
approaching 2, i.e., the discrepancy between the RRL and
Commander solution is worse for very extended emission. This
is likely to be due to excess synchrotron radiation at higher fre-
quencies that has been accounted for in the simple Commander
model by the free-free component. In particular, the synchrotron
component is effectively modelled as a power law in frequency
with a fixed spectral index, while there is considerable evi-
dence for spectral flattening of synchrotron emission at low lat-
itudes (Kogut et al. 2007;de Oliveira-Costa et al. 2008;Gold
et al. 2011;Planck Collaboration Int. XXIII 2015). This will
contribute to some level of excess emission at frequencies 20–
100 GHz, which can result in an apparent increase in the free-
free and/or AME amplitude in the Commander fits, an effect that
can be seen as a broad background (zero-level) in the longi-
tude plots of Fig. 9. We also note that the RRL data will not
reliably trace large-scale (&4◦) emission, since the observations
were made in 8◦×8◦patches of sky.
In summary, the Commander free-free map appears to be a
reliable tracer of the brightest H ii regions, with an accuracy
of around 20%. Weaker emission regions appears to be over-
estimated by up to a factor of 2, relative to estimates using RRL
data. As mentioned in Planck Collaboration X (2016), including
new data in the frequency range 2–20 GHz will yield improved
component separation products, particularly for synchrotron and
AME, which in turn will improve the accuracy of the free-free
solution.
4.2. Anomalous microwave emission
Commander models the AME as spinning dust, which for uni-
form conditions (grain size distribution, ambient radiation) gives
a relatively sharply peaked spectrum; specifically it uses the
SpDust2 code (Ali-Haïmoud et al. 2009;Silsbee et al. 2011) to
calculate a template spectrum using parameters typical for the
diffuse cold neutral medium (Draine & Lazarian 1998). To toler-
able accuracy for present purposes, models for other phases can
be approximated by scaling the template in frequency.
As shown in Sect. 3, the AME spectrum is quite vari-
able, with an apparent tendency to peak at substantially higher-
than-average frequencies around some H ii regions. Low-latitude
lines-of-sight will therefore contain AME with a range of peak
frequencies, which will give a broader spectrum that cannot be
fitted by a frequency-shifted template. In fact, a superposition
of AME spectra with a distribution of peak frequencies could
approximate a power law over the frequency range of primary
interest here (20–70 GHz). Such a general model could not be
constrained by the available data, but we take a step in this di-
rection by using two AME components in the Commander fit,
each based on our SpDust2 template. One component (which
generally dominates) is scaled in frequency at each pixel, with
a Gaussian prior for the frequency of the intensity peak, νp1,
of (19 ±3) GHz. The peak frequency of the other component
A25, page 10 of 45
Planck Collaboration: Planck 2015 results. XXV.
Fig. 8. Scatter plots of the Commander versus the RRL free-free map for three regions of the Galactic plane (the first three regions shown in Fig. 7).
The slope values given in each panel result from a linear fit to all the data points.
50 40 30 20 10 0
Longitude [degree]
0
5
10
15
20
25
Temperature [K
RJ
]
350 340 330 320 310
Longitude [degree]
0
5
10
15
20
25
Temperature [K
RJ
]
300 290 280 270 260
Longitude [degree]
0
5
10
15
Temperature [K
RJ
]
Fig. 9. Longitude profiles of the Commander (black) and RRL (red) free-free maps, averaged within |b| ≤ 1◦, for three regions of the Galactic plane
(the first three regions shown in Fig. 7). The comparison is for the bright compact regions that show as peaks on top of a broad baseline level due
to extended emission.
is assumed to be constant over the whole sky with a value to
be determined globally; in this Commander solution the result is
νp2 =33.35 GHz. This second model component is there to ac-
count for the broadened spectrum of the overall AME emission,
and should not be interpreted as a physically distinct AME com-
ponent; in this paper (as in Planck Collaboration X 2016) we
only consider the combination of the two components, which we
evaluate at 22.8 GHz to allow comparison with earlier studies.
The all-sky map of AME as fitted by Commander is shown in
Fig. 10. The foreground emission detected at frequencies around
20–60 GHz in the WMAP/Planck maps appears to be domi-
nated by AME (e.g., Davies et al. 2006;Ghosh et al. 2012;
Sect. 3of this paper). It has been detected with high signif-
icance over a large area of the sky. The strongest emission is
observed in the Galactic plane, particularly in the inner Galaxy
(l=300◦→0◦→60◦). The AME in the inner plane was iden-
tified and discussed by Planck Collaboration XXI (2011) and
Planck Collaboration Int. XXIII (2015). The map also shows
clear, strong AME in many of the molecular cloud regions
that are well known to exhibit AME, including the Perseus
and ρOphiuchus clouds. While previous Planck AME papers
(Planck Collaboration XX 2011;Planck Collaboration Int. XV
2014) focused on compact AME regions, here we also look at
diffuse AME regions away from the Galactic plane (|b|>10◦).
We identify strong diffuse AME in the Chameleon/Musca re-
gions as well as the Orion complex, which is discussed later in
this section. These regions are less likely to be affected by fore-
ground degeneracies than others: regions such as the Auriga dust
feature above the Galactic centre and a large structure around the
North Celestial pole can also be seen in the AME map, however
they overlap with significant sources of synchrotron emission.
Overall there is a high level of correlation of AME with the
thermal dust maps made at higher frequencies of Planck and in
the infrared; however, the AME clearly cannot be accounted for
by the Rayleigh-Jeans tail of the thermal dust emission, which
is at least a factor of 30 below the AME component at 30 GHz.
At low Galactic latitudes (|b|.10◦) and in bright compact and
diffuse AME regions, the complexity of separating AME from
CMB, synchrotron, and free-free emission is the dominant un-
certainty for the Commander solution. At high latitude, where the
AME is weak, the Commander solution is dominated by instru-
mental noise. The Commander AME map also contains residual
extragalactic source contributions when their spectrum is inter-
mediate between the assumed synchrotron and free-free models
(e.g., Centaurus A is clearly visible, as are bright radio sources
such as 3C 273); to remove these we mask sources in the PCCS2
catalogue at 28.4 GHz that are brighter than 1 Jy.
We evaluate the peak frequency of the combined AME spec-
trum for each pixel. The mean and median across the whole
sky are 20.6 and 20.4 GHz, respectively, and at |b|>10◦they
are 20.5 and 20.2 GHz, respectively. Apart from bright regions,
the peak frequency map is very noisy, largely constrained by
our prior on νp1 around 19 GHz (see Sect. 4.1 and Fig. 16 of
Planck Collaboration X 2016) and contains residual signal from
point sources. Regions such as Perseus, ρOph and Orion have
higher than average peak frequencies, around 25–30 GHz, in
agreement with previous analyses (e.g., Planck Collaboration
XX 2011). In some places the fits are dominated by the sec-
ond AME component (notably the Sh 2-27 and southern Gum
nebulae). This is consistent with the ILC analysis of Sect. 3, al-
though in reality some of these regions have even higher peak
frequency, e.g., Planck Collaboration Int. XV (2014) found a
A25, page 11 of 45
A&A 594, A25 (2016)
Fig. 10. All-sky map of AME from Commander at 22.8 GHz plotted using a Mollweide projection, with a 30◦graticule and an asinh colour scheme.
Seven regions of diffuse AME have been highlighted, and are discussed in the text.
Fig. 11. Mask used for the AME T–Tplots. The regions in grey are
masked in the T–Tplots. The region in light blue shows where the AME
at 22.8 GHz is brighter than 1 mK; this is also masked when determining
the best-fitting emissivity. We also mask the ecliptic plane (|β|<20◦)
when analysing the IRAS and WISE data.
best-fitting value of νpof 50 GHz for the California Nebula. Sh2-
27 has an exceptionally high χ2in the Commander fit (Planck
Collaboration X 2016, Fig. 22), which we found can be alle-
viated by allowing a similar peak frequency (recall that in our
baseline Commander model νp2 is fitted globally to keep the
number of free parameters acceptable). Dobler & Finkbeiner
(2008b) previously found a bump at ∼50 GHz in the SED of
Hα-correlated emission from the Warm Ionized Medium that
they speculated could be a high-frequency spinning dust compo-
nent, which is likely the same emission we are seeing here. We
find that in bright free-free regions (primarily on the plane, but
also bright high-latitude regions such as M 42 and the California
nebula) the AME peak frequency correlates with the free-free
amplitude (at 22.8 GHz) with a best-fitting slope of (0.066 ±
0.012) GHz mK−1, which may be an indication of the importance
of the interstellar radiation field in the excitation of AME carriers
(e.g., Tibbs et al. 2013).
4.2.1. All-sky correlations
In order to look at the correlation of the AME with other com-
ponents, we have degraded the AME map and other tracer maps
to Nside =64 (pixel widths of 550). We define a mask, shown in
Fig. 11, to exclude pixels where the S/N for AME is less than
3.0 and point sources brighter than 1 Jy, except at |b|<5◦to
avoid masking Galactic plane sources. However, we mask five
Galactic plane regions containing the brightest H ii complexes
(e.g., Cygnus-X) in the sky. We use a 3◦radius disc centred
at (l,b)=(49.
◦6,−0.
◦6), (79.
◦2,0.
◦4), (268.
◦2,−1.
◦3), (287.
◦5,−1.
◦1)
and (291.
◦8,−0.
◦9). These are particularly bright in the IRAS 12
and 25 µm maps, where there could be contamination by starlight
and/or line emission. We show the T–Tplots for the entire sky
outside of the mask for different potential tracers in Fig. 12.
We also test the effects of additionally masking the Galactic
plane (b<10◦). We calculate the emissivity, i.e., the ratio of
AME emission compared to AME tracers, from the slope of
the best-fitting line as shown in the T–Tplots for pixels where
the AME amplitude at 22.8 GHz is less than 1 mK (shown in
light blue in Fig. 11). The difference between the Commander
maximum-likelihood and mean solutions for the AME com-
ponent is around 5–6%, depending on frequency, due to the
non-Gaussian shape of the probability distribution. We there-
fore assume a conservative 10% modelling uncertainty on top
of the Commander component uncertainty maps; we also as-
sume a 13.5% calibration uncertainty for the IRAS 100 µm map
(Miville-Deschênes & Lagache 2005).
We find an excellent correlation (r=0.98) with the Planck
map at 545 GHz, shown in Fig. 12a, which is predominantly ther-
mal dust emission. A mean emissivity of (65±7) µK (MJy sr−1)−1
is found ((70 ±7) µK (MJy sr−1)−1when the Galactic plane
A25, page 12 of 45
Planck Collaboration: Planck 2015 results. XXV.
Fig. 12. T–Tplots at Nside =64, using the mask shown in Fig. 11. The best-fitting lines and emissivities shown are determined where the AME
amplitude is less than 1 mK. The Spearman’s rank correlation coefficient (s-rank) and the Pearson’s correlation coefficient (r) are also shown. Top
row: AME (µK at 22.8 GHz) compared to: a) Planck 545 GHz (MJy sr−1); b) the Commander thermal dust solution at 353 GHz (µK); and c) τ353
(rescaled by 105). Middle row:d) 100 µm (MJy sr−1); e) Hi(K km s−1); and f) CO J=2→1 (K km s−1). Bottom row:g) IRAS 12 µm (MJy sr−1)
divided by G0;h) IRAS 25 µm (MJy sr−1) divided by G0; and i) dust radiance calculated from the Commander products.
is masked), which is consistent over a wide range of ampli-
tudes. The overall emissivity against IRAS 100µm that we
find when looking at the entire sky is (22 ±2) µK (MJy sr−1)−1
(Fig. 12d), with r=0.91; this is somewhat lower than the
weighted average found by Planck Collaboration Int. XV (2014)
of (32 ±4) µK (MJy sr−1)−1, although the latter is for a specific
set of 27 low-latitude clouds selected for significant emission,
rather than an all-sky average. Davies et al. (2006) used tem-
plate fitting to estimate the emissivity of AME compared with
thermal dust; they calculated the emissivity in terms of the “FDS
model 8” dust map at 94 GHz (Finkbeiner et al. 1999), which
can be multiplied by a factor of 3.3 to obtain the equivalent
100 µm emissivity (on average). Our analysis agrees well their
results, where they found (21.8±1.0) µK (MJy sr−1)−1across the
whole sky outside of the WMAP Kp2 mask, and an average of
(25.7±1.3) µK (MJy sr−1)−1in the 15 regions they examined. We
note that this good agreement in the overall high-latitude emis-
sivity comes from different component-separation methods.
For thermal emission tracers there is some curvature in the
relation at the brightest pixels, where there is less AME for a
given amount of thermal dust emission. This is more notice-
able in the 100 µm correlation than the 545 GHz correlation, be-
cause the 100 µm amplitude is much more affected by the tem-
perature of the thermal dust grains than the 545 GHz emission
(Tibbs et al. 2012), which increases towards the Galactic plane
(e.g., Planck Collaboration XIX 2011). This effect is not seen in
the τ353 correlation, however there is a much larger scatter in
the τ353 correlation, as well as a trend towards there being more
AME per unit τ353 at high optical depth.
There is a good correlation (r=0.84) with H iemission inte-
grated in the velocity range [−450,+400] km s−1from the LAB
survey (Kalberla et al. 2005), shown in Fig. 12e, with a slope of
(210 ±20) µK K−1((280 ±30) µK K−1when the Galactic plane
is masked); there is also a noticeable upturn and increase in the
scatter of points in this relation above an H iintegrated velocity
intensity of around 3 K, in agreement with previous results (e.g.,
A25, page 13 of 45
A&A 594, A25 (2016)
Lagache 2003). Planck Collaboration Int. XVII (2014) com-
pared Planck and WMAP data to H idata, and found an
emissivity of (1.75 ±0.14) ×107µKτ−1
353
7; they also con-
sidered AME models, which gave slightly lower emissivities
with τ353 of 1.3×107µK. We find an emissivity against τ353
of (8.7±0.8) ×106µK (Fig. 12c, (1.1±0.1) ×107µK ex-
cluding the Galactic plane), which is lower than the values
found by Planck Collaboration Int. XVII (2014) by a factor of
1.5–2.1. However, the emissivities have not been calculated
on the same areas of sky, and this difference is within the
variations of E(B−V)/NHreported in Planck Collaboration
XI (2014). We also compare AME at 22.8 GHz with the
Commander thermal dust solution at 353 GHz and find a
slope from the regions of (0.9±0.1) K K−1(Fig. 12b, r=0.98,
(1.0±0.1) K K−1excluding the Galactic plane); this is also
slightly lower than the value of (1.12 ±0.03) K K−1found by
Planck Collaboration Int. XXII (2015).
The upturn in the H icorrelation implies that the AME is
related to another phase that is not traced by H iemission.
We also look at CO emission, which will trace denser regions,
and is a marker for H ii emission. The correlation with the
Commander CO J=2→1 map is rather more complicated
than the other correlations, however. The best-fitting slope is
(180 ±20) µK (K km s−1)−1(Fig. 12f), with r=0.91, reducing to
(150 ±20) µK (K km s−1)−1when the Galactic plane is masked.
However, at low amplitudes the emissivity is clearly a lot lower
than the best-fitting value, and at high amplitudes it is clearly
higher than the best-fitting value. This could be due to compo-
nent separation issues; alternatively, since the curvature largely
takes place at AME amplitudes above 1 mK, it may indicate that
there is less AME in the densest parts of the Galaxy than at
higher latitudes.
We have also looked at the correlations with the IRAS 12 µm
and 25 µm data, which trace the very small dust grain popula-
tion that could produce AME, as well as containing PAH line
emission. Indications of an improved correlation with these data
have been seen on smaller scales by e.g., Casassus et al. (2006).
However, these data are considerably contaminated by zodiacal
light and other systematics; as such, we have masked the ecliptic
plane (|β|<20◦) before fitting these data. The resulting AME
emissivity against 12 µm emission is (460 ±60) µK (MJy sr−1)−1,
with Pearson correlation coefficient r=0.93 and Spearman’s
rank correlation coeffient s=0.71, and against 25 µm emis-
sion is (370 ±40) µK (MJy sr−1)−1, with r=0.87 and s=0.52.
However, these will still include other systematic effects, such
as stellar contamination and high-latitude zodiacal light. We
nomalize the maps by the interstellar radiation field intensity,
which we estimate by G0=(Td/17.5)β+4(Ysard et al. 2010),
so that we are correlating against column density. Dividing
by G0improves the correlation as measured by Spearman’s
rank, which includes the curvature in the correlation, while the
Pearson’s rank, which relies on a linear correlation, reduces.
This improvement in correlation is as expected by the spin-
ning dust model (Ysard et al. 2010), however the value of G0,
and hence the correlation coefficient, is highly dependent on the
dust temperatures and spectral indices. The resulting emissivity
against 12 µm is (290 ±30) µK (MJy sr−1)−1(Fig. 12g, r=0.78,
s=0.87) and against 25 µm it is (200 ±20) µK (MJy sr−1)−1
7Planck Collaboration Int. XVII (2014) give a value of
(17 ±1) µK (1020 H cm−2)−1at 23 GHz, which can be converted to
µKτ−1
353 using E(B−V)/NH=(1.44 ±0.02) ×10−22 mag cm2and
E(B−V)/τ353 =(1.49 ±0.03) ×104mag−1cm−2(Planck Collaboration
XI 2014).
(Fig. 12h, r=0.77, s=0.83). We return to this topic in the dis-
cussion of the Musca region, below.
In a recent analysis, Hensley & Draine (2015) have com-
pared the Commander AME maps with different potential tracers
of AME, and find a good correlation against dust radiance. This
is the integrated intensity from thermal dust, and it characterises
the energy absorbed and emitted by the thermal dust particles.
We have also compared the AME map with dust radiance (cal-
culated using the Commander products and Eq. (10) in Planck
Collaboration XI 2014) in Fig. 12i. We find a strong correlation
at low amplitudes, with an emissivity against AME at 22.8GHz
of (5.6±0.6) ×108µK, and r=0.90. At higher amplitudes (i.e.,
in the Galactic plane), the dust radiance calculated using the
Commander products turns over, in the same way as the corre-
lation with the IRAS 100 µm data does, implying that this is due
to temperature effects. If we use the radiance map from Planck
Collaboration XI (2014), which includes the 100 µm map, in the
analysis, this effect reduces in magnitude, but is still present.
Hensley & Draine (2015) also compare the variations in
AME emissivities (defined by dividing the AME map by the
average emissivity against dust radiance) with variations in the
fraction of the WISE 12 µm data (Meisner & Finkbeiner 2014)
that can be attributed to polycyclic aromatic hydrocarbon (PAH)
emission, in contradiction with expectations from the spinning
dust model where PAHs are though to be the spinning molecules.
Using the Commander thermal dust amplitude and Tdand βd
maps from (Planck Collaboration XI 2014), they find no cor-
relation between AME emissivities and PAH emission. We re-
peat their analysis, using the WISE 12 µm data and masking the
Galactic plane (|b|<5◦) and the Ecliptic plane (|β|<20◦) in ad-
dition to the mask shown in Fig. 11. We find similar results when
using the dust radiance maps calculated using the products from
(Planck Collaboration XI 2014), with correlation coefficients of
0.10 (Planck Collaboration XI 2014 radiance map) and 0.23 (us-
ing the Commander dust amplitude). However, if we calculate
dust radiance using only the Commander products, then we find
a correlation with a slope of 1.1±0.1, with a correlation coeffi-
cient of 0.52 8. Although the Commander dust radiance map will
be biased low due to the absence of data points tracing warmer
dust temperatures, this demonstrates the dependence of this re-
sult on the quality of the radiance map; it will also depend on
the quality of the AME map. One possibility is that this correla-
tion is affected by the large angular scale structure in the Galaxy;
as such we return to this in the discussion of the Musca region,
below.
In conclusion, we find the best correlation with AME at all
amplitudes is from the 545 GHz Planck map, followed by the
Commander dust solution at 353 GHz and the optical depth, τ353.
The dust radiance has a tight correlation with AME away the
Galactic plane, but has a worse correlation in the plane. The cor-
relation with 100 µm is significantly affected by temperature ef-
fects, which the choice of lower frequency thermal dust maps
avoids. We find a reasonable correlation with H iemission, al-
though this does not appear to correlate with the brightest AME
emission, and with CO emission, although this is not well-fitted
with a single emissivity. The correlation with dust radiance is
very good at low amplitudes, but using a single emissivity would
8The first preprint version of Hensley & Draine (2015) used a full-sky
WISE map that later turned out to be contaminated by Planck 857 GHz
data on scales larger than 2◦, which might also affect our analysis of
their results. A revised version, Hensley et al. (2016), used an alternative
approach to process the WISE data to avoid this issue; they found that
their conclusions were unchanged. As such, we have not revised our
conclusion here.
A25, page 14 of 45
Planck Collaboration: Planck 2015 results. XXV.
Table 2. Top section: emissivities of AME at 22.8 GHz (µK) relative
to the Commander thermal dust amplitude at 545 GHz (MJy sr−1), the
IRAS map at 100 µm (MJy sr−1), and the optical depth at 353 GHz,
τ353, for the regions shown in Figs. 10 and 13.Bottom part: emissivities
from Davies et al. (2006, D06; whole sky and region mean) and Planck
Collaboration Int. XV (2014, XV; Perseus, ρOph; and the unweighted
region mean) for comparison.
Region . . . . . . . . AME/545 GHz AME/100 µm AME/τ353
[µK (MJy sr−1)−1] [µK (MJy sr−1)−1] [µK 10−6]
R1: Perseus . . . . . 24 ±7 12.3±1.9 1.5±0.9
R2: Plume . . . . . 47 ±6 18 ±2 7.7±1.0
R3: RCrA . . . . . 36 ±14 50 ±12 4.1±1.8
R4: ρOph . . . . . . 40 ±9 4.6±0.9 2.2±1.2
R5: Musca . . . . . 59 ±8 26 ±3 6.9±1.0
Chamaeleon . 74 ±8 22 ±2 11 ±1.1
R6: Orion . . . . . . 47 ±5 20 ±2 4.7±0.6
R7: λOrionis . . . 104 ±11 25 ±3 15 ±1.8
LMC . . . . . . . . . 56 ±6 8.5±1.0 7.9±0.9
SMC . . . . . . . . . 30 ±9 4.7±1.5 3.3±1.0
Entire sky . . . . . . 65 ±7 22 ±2 8.3±0.8
|b|>10◦. . . . . . . 70 ±7 21 ±2 9.7±1.0
XV: Perseus . . . . . . . 24 ±4. . .
XV: ρOph . . . . . . . . 8.3±1.1. . .
XV: Mean . . . . . . . . . 32 ±4. . .
D06: Kp2 mask . . . . . 21.8±1.0. . .
D06: Region mean . . . 25.7±1.3. . .
Notes. We also include emissivities for the LMC and SMC, discussed
in Sect. 4.5.
over-predict the amount of AME present in the Galactic plane.
We caution that these correlations will depend on the choice of
mask due to large-scale biases, contamination from other emis-
sion mechanisms (e.g., point sources, zodiacal light, etc.), and
variations across the sky.
4.2.2. Diffuse AME regions
We now move on to consider individual diffuse regions of AME.
We have labelled seven regions in the all-sky AME map of
Fig. 10 that demonstrate diffuse AME away from the Galactic
plane and are in areas with relatively high S/N and clean compo-
nent separation outputs. These include the well-known regions
of Perseus and ρOphiuchus, as well as five new regions. In
Fig. 13 we show the Commander solutions for synchrotron, free-
free, AME and thermal dust, along with the CMB-subtracted
28.4 GHz Planck data and Hαdata, for the new regions. In each
of these maps, an area has been defined to analyse the properties
of the region using T–Tplots, and sources brighter than 5 Jy in
the PCCS2 catalogue have been masked out to a radius of 600,
with some exceptions noted below. We show T–Tplots between
AME and thermal dust for three regions in Fig. 14. We use the
best-fitting slope from these T–Tplots to determine the emissiv-
ity of the AME component compared to the other components.
Results for all regions are given in Table 2.
1. Perseus. This molecular cloud is a well-known source of
AME (Watson et al. 2005;Planck Collaboration XX 2011). For
the correlation analysis, we focus on an 8◦×8◦patch cen-
tred on (l,b)=(160.
◦26,−18.
◦62). Perseus is included in
the PCCS2 catalogue, so we unmask the point sources when
calculating the emissivities in this region. We find a lower
emissivity for Perseus in the Commander products than in
Planck Collaboration Int. XV (2014); this is due to leakage of
around 40% of the AME emission to the free-free component,
which also causes a decrement in the synchrotron map (see
Fig. 33 below).
2. Pegasus plume. An example where the AME amplitude
traces a filament-like structure at the edge of the Pegasus con-
stellation at (l,b)=(92◦,−37◦) and approximately 10◦in length
is shown in the first column of Fig. 13 (Region 2 in Fig. 10).
We look at a 15◦×15◦patch centred on this position. The
plume was included in the CO catalogue of Magnani et al.
(1985) as MBM 53–55 (G92.
◦97−32.
◦15, G92.
◦97−37.
◦54, and
G89.
◦19−40.
◦94, respectively), and MBM 55 is also coincident
with the H ii region S 122 (G89.
◦18−41.
◦13) in Sharpless (1959).
Fukui et al. (2014) have compared the CO emission to thermal
dust emission as seen by Planck at 353–857 GHz. The struc-
ture can be seen clearly in the CMB-subtracted Planck 28.4 GHz
map, which is very closely correlated with the thermal dust am-
plitude at 545 GHz. There is no visible emission in the syn-
chrotron map and thus the signal in the AME map is unlikely
to be due to synchrotron emission. The feature does have associ-
ated free-free emission, as traced by Hα; however, the brightness
at 22.8 GHz is about 10 times greater than would be predicted
for free-free emission, assuming a typical electron temperature
(Te≈7000 K). The plume structure is not resolved at 1◦resolu-
tion, and therefore absorption of Hαlight by dust should be sim-
ilar to other high-latitude regions (|b|>30◦) at .0.1 mag, with
only a small fraction being absorbed in compact high-density
regions. The section of the plume corresponding to S 122 is def-
initely free-free emission; however, we conclude that the rest of
this feature is predominantly AME, and that its appearance in the
Commander free-free map is likely to be due to leakage from the
AME component. This would be an interesting target for higher-
resolution follow-up measurements, particularly in the bright re-
gion around MBM 53.
3. Corona Australis. The second column of Fig. 13 shows
Corona Australis below the Galactic centre (Region 3 in Fig. 10).
The region has been studied by Harju et al. (1993) in CO emis-
sion and with Spitzer by Peterson et al. (2011), and it is re-
viewed in Neuhäuser & Forbrich (2008). The central object is
the R Corona Australis dark cloud (R CrA), which has a tail ex-
tending to the bottom-left of the maps consisting of the reflection
nebula NGC 6729. The reflection nebula is illuminated by the
early-type star TY CrA, and exhibits an extended 3.3 µm emis-
sion feature that could be to PAH emission (Chen & Graham
1993). We focus our analysis on a 3◦×3◦area centred on
(l,b)=(1◦,−22◦) that encompasses the reflection nebula but
not the R CrA cloud. The R CrA region is dominated by free-
free emission, and it appears in both the Hαand free-free maps.
However, the reflection nebula clearly shows AME, with a small
amount of free-free emission seen in Hα(there is a feature in the
tail in the Commander free-free map; this does not show up in
Hα, however). The structure does not appear in the synchrotron
map.
4. ρOphiuchus. This molecular cloud is another
well-known source of AME (Casassus et al. 2008;
Planck Collaboration XX 2011). It is included in the
PCCS2 catalogue, so we unmask the point sources
when considering this region. We look at a 5◦×5◦
patch centred on (l,b)=(353.
◦05,16.
◦9). We find a low
A25, page 15 of 45
A&A 594, A25 (2016)
Fig. 13. Commander component maps for diffuse AME regions. From top to bottom: synchrotron amplitude (mK at 22.8 GHz), Hα(Dickinson
et al. 2003, Rayleighs), free-free amplitude (mK at 22.8 GHz), AME amplitude (mK at 22.8GHz), thermal dust amplitude (MJy sr−1at 545 GHz),
CMB-subtracted Planck 28.4 GHz map (mK). Column 1: a plume of emission in Pegasus, centred at (l,b)=(91.
◦5,−35.
◦8). Column 2: the Corona
Australis region at (l,b)=(0◦,−18◦). Column 3: an extended region of emission in Musca/Chamaeleon, centred at (l,b)=(305◦,−26◦) (the Large
Magellanic Cloud can be seen to the bottom-right). Column 4: the Orion (M42) region of emission centred at (l,b)=(209◦,−19.
◦38). Column 5:
the λOrionis region at (l,b)=(196◦,−12◦). The graticule separation is 10◦in both directions, and the colour scales are asinh (although most are
close to linear). The regions looked at for the correlation analyses are shown as black rectangular boxes.
A25, page 16 of 45
Planck Collaboration: Planck 2015 results. XXV.
Fig. 14. T–Tplots comparing the Commander AME map evaluated at 22.8GHz to the Commander dust map at 545 GHz in three regions, as shown
in Fig. 13. The best fit is shown as a dashed line. From left to right, for the Pegasus plume, Musca, and λOrionis regions.
AME-100 µm emissivity of (4.6±0.9) µK (MJy sr−1)−1,
compared to (8.3±1.1) µK (MJy sr−1)−1from Planck
Collaboration Int. XV (2014). This is due to significant leakage
between the AME, free-free, and synchrotron components in
the Commander solution in this region: half of the emission
that has been attributed to AME in previous analyses (e.g.,
Planck Collaboration XX 2011) is instead attributed to free-free,
causing a decrement in the synchrotron map, similar to the
Perseus region (see Fig. 33 below). This issue is due to the
complexity of this region, which lies within the Gould belt,
where there is bright surrounding free-free and synchrotron
emission.
5. Musca/Chamaeleon region. The third column of Fig. 13
shows a region of diffuse AME below the Galactic plane, in
the Musca region close to the Magellanic clouds (Region 5
in Fig. 10). This region was previously studied in dust polar-
ization by Planck Collaboration Int. XX (2015). Its distance is
around 160–180 pc, and it contains three dark clouds (Whittet
et al. 1997). We focus on a 15◦×12◦region centred on (l,b)=
(299◦,−16.
◦5). The AME is clearly correlated with the dust emis-
sion, and there is a notable absence of emission in the syn-
chrotron and Hαmaps in this region. There is structure evident
in the Commander free-free map; however, this correlates with
some parts of the dust emission and not Hα, so it is likely to
be AME rather than free-free. We mask two bright thermal dust
sources in the Musca region (Cha I and Cha II) out to a radius of
800when calculating the emissivities.
There is also a much larger AME half-ring that extends south
of the Musca feature, into the Chamaeleon constellation. We
include emissivities for the Chamaeleon region in a 35◦×30◦
region centred on (l,b)=(305◦,−25◦). We also use this re-
gion to look at the correlations between AME and IRAS 12 and
25 µm, since this part of the sky is at both high Galactic and
ecliptic latitudes, and as such it is less contaminated by zodia-
cal or stellar light than the rest of the sky. We find an emissivity
against 12 µm of (550 ±60) µK (MJy sr−1)−1, and against 25 µm
of (570 ±70) µK (MJy sr−1)−1, both of which are higher than the
all-sky values given earlier in this section. We also divide the 12
and 25 µm data by G0(as described above), after which we find
an emissivity against 12 µm of (250 ±25) µK (MJy sr−1)−1, and
against 25 µm of (159 ±16) µK (MJy sr−1)−1. That these num-
bers are lower than for the whole sky is likely to be an arte-
fact in the G0map, which has a very small value due to the low
Commander dust temperature in this region of around 17 K, com-
pared to the mean temperature of 20.9 K. We also correlate the
variations in AME emissivities with the fraction of PAH emis-
sion at 12 µm (as described above for the whole sky, following
Hensley & Draine 2015) in this region. We find similar results to
above, with a slope of 3.4±0.4 when using the Commander ra-
diance map, but poor correlations when using the radiance maps
based on Planck Collaboration XI (2014). The correlations pre-
sented here are sensitive on how the comparison maps have been
made, both in terms of inputs to the radiance map and contami-
nating emission in the IRAS maps.
6. Orion region. The fourth column of Fig. 13 shows the
Orion region (Region 6 in Fig. 10). The strong free-free emis-
sion from M 42 and Barnard’s Loop (Sh 2-276) is clearly visible.
However, there is also an arc of AME emission that correlates
with the high-frequency dust emission. This arc extends to lower
longitude from M 42 along the integral-shaped filament to the
dark cloud L 1641 and perpendicularly crosses over Barnard’s
loop (Orion Molecular Cloud A); in the other direction (Orion
Molecular Cloud B) it extends up to M 78 to the top-right of
the figure. The arc also appears faintly in free-free emission
and in Hα. This dusty feature was recently mapped in 3D by
Schlafly et al. (2015). We look at a 15◦×15◦region centred on
(l,b)=(209.
◦01,−19.
◦38), and we additionally mask M 42 to a
radius of 1000, and the point where the dust crosses the top part
of Barnard’s loop (where AME has likely leaked into the free-
free map), when calculating the emissivities.
7. λOrionis. The λOrionis region (Region 7 in Fig. 10) can
also be seen to the right of Barnard’s Loop, and is shown in the
fifth column of Fig. 13. We focus on a 13◦×13◦region centred on
(l,b)=(195.
◦2,−12.
◦2). This H ii region (also known as S 264),
which is illuminated by the O8 star λOrionis, exhibits a shell
of AME around the outside of the free-free region coincidental
with the thermal dust ring. The dust ring has been previously
noted in Maddalena & Morris (1987) and Zhang et al. (1989),
and it has been observed in H iby Wade (1957) and CO by Lang
& Masheder (1998) and Lang et al. (2000). The brightest AME
regions in the ring are mostly located adjacent to bright ther-
mal dust features, although the brightest AME region does not
have a bright thermal dust counterpart. AME has been seen in
PDRs around other H ii regions at higher resolution, e.g., in ρ
Ophiuchus (Casassus et al. 2008) and Perseus (Tibbs et al. 2010).
λOrionis has a particularly high emissivity against 545 GHz and
τ353, but the emissivity against 100 µm is comparable to the av-
erage; this indicates that the AME is connected to the colder dust
in this region. It is likely that this region has not been identified
A25, page 17 of 45
A&A 594, A25 (2016)
before due to the large free-free emission feature that the ring of
AME surrounds.
There are significant differences between the emissivities for
these seven regions. Some emissivities (particularly for Perseus
and ρOphiuchus) are clearly biased low due to component-
separation issues. Four are relatively consistent with values of
AME/545 GHz typically in the range 40–50 µK (MJy sr−1)−1.
The notable exceptions are the emissivity values for λOrionis
and the Chamaeleon region, which are higher than the other
regions by a factor of 2. However, except for λOrionis and
Chamaeleon, the emissivities are significantly lower than the av-
erage emissivity of (65 ±7) µK (MJy sr−1)−1across the whole
high-latitude sky. The differences in emissivities could be a
component-separation artefact, or they could be an indicator of
the dependence of environmental conditions for the AME, or
dust grain size distribution. For example, in nearby molecular
clouds the AME carrier grains could have attached themselves
to larger dust grains more than in the diffuse medium (e.g., Kim
et al. 1994). We find a good correlation between the emissivities
at 545 GHz and the optical depth.
In conclusion, we find that the Commander AME map pro-
vides a good tracer of AME in our Galaxy, however, there are
significant degeneracies between the free-free and AME compo-
nents that present difficulties when using the map to calculate
emissivities. Additional all-sky data at frequencies of 5–20 GHz
are needed to improve the free-free tracer and so enable a cleaner
separation of AME from the other low-frequency components,
in order to determine accurate emissivities and comparison with
other data sets. In particular, λOrionis is a distinctive region that
warrants further study.
4.3. Synchrotron
In this section we discuss constraints on synchrotron emission
derived from total intensity. We discuss options for modelling the
form of the synchrotron spectrum, which is not a simple power
law, and specify the model used in our baseline Commander anal-
ysis. The results are strongly limited by uncertainties in compo-
nent separation; a much more reliable picture of the structure
of the high-frequency synchrotron emission emerges from po-
larization data, discussed in Sect. 5. However, one advantage of
total intensity is that we have a high S/N data set in the 408 MHz
Haslam map, whereas in polarization all available ground-based
sky maps are at low frequency and hence too strongly affected
by Faraday rotation to be useful for spectral analysis. Because
of residual differences between Planck and WMAP polarization
maps (see Sect. 5.1 and Planck Collaboration X 2016), we do
not attempt an independent fit of the spectrum of the polarized
synchrotron emission.
The Galactic synchrotron spectrum curves significantly be-
low a few GHz (e.g., de Oliveira-Costa et al. 2008;Strong et al.
2011). To generate a useful spectral constraint from just one
low-frequency map, we cannot afford to fit this curvature inde-
pendently at each pixel, as otherwise the low-frequency point
would always fit perfectly and give no constraint on the high-
frequency spectrum; instead we force the curvature to be con-
stant across the sky, and we have tried several approaches to
regularize the fit. As expected, simple power-law fits are incon-
sistent with our data, generating large and spurious “gain correc-
tions” at 408 MHz.
Strong et al. (2011) and Orlando & Strong (2013) model
the observed synchrotron emission for given Galactic magnetic
Fig. 15. Local spectral index of the synchrotron emission β(ν)=
dln T/dln νvs. frequency for a sample of pixels (one per Nside =8
super-pixel), in the GALPROP z10LMPD_SUNfE model from Orlando &
Strong (2013). The spectra are colour-coded by Galactic latitude: spec-
tra at low latitudes show strong low-frequency curvature due to free-free
absorption.
fields and cosmic-ray (CR) propagation scenarios with the
GALPROP code (Strong et al. 2007). This code solves the CR
transport equation for any CR species, accounting for diffusion,
reaccelerating processes, and energy losses in the interstellar
medium. The CR transport properties are constrained by the lo-
cal CR measurements and the observed ratio of secondary to
primary nuclei, while energy losses are calculated for a given
Galactic magnetic field. As implemented in the papers above,
synchrotron GALPROP models aim to reproduce the large-scale
emission, but not sub-kpc scale features such as the synchrotron
loops and individual supernova remnants. The modelled local in-
terstellar energy spectrum of leptons (after propagation effects)
is adjusted to reproduce the direct measurements by Fermi-LAT
above 7 GeV (Abdo et al. 2009;Ackermann et al. 2010). At
lower energies, where solar modulation makes the local spec-
trum hard to determine, the lepton spectrum is based on syn-
chrotron observations. In order to reproduce the curvature of the
spectrum below a few GHz, Strong et al. (2011) found that the
injected electron spectrum (before propagation effects) should
have a break at approximately 4 GeV and should be harder than
p=1.6 below 4 GeV. Accounting for all these constraints, the
resulting injected spectral index used in those works are (N(E)∝
E−p) with p=1.6/2.5/2.2, respectively, below/between/above
the breaks at 4 and 50 GeV. Moreover, the cosmic ray lepton
(CRL) spectral shape varies with Galactocentric radius due to
propagation effects, CR source distribution, and the magnetic
field strength. Despite these effects, the analysis in Strong et al.
(2011) and Orlando & Strong (2013) finds that the spectral index
at any given frequency is strikingly uniform across the sky: the
spectral index between 408 MHz and 22 GHz has an rms disper-
sion of 0.02; between 22 and 70 GHz the rms is 0.01 (Fig. 15).
The uniformity of these model spectra over the sky contrasts
with the much larger variation in β(ν) with frequency in the same
models. For a given CRL spectrum, the frequency of synchrotron
radiation scales directly with B⊥, the component of the magnetic
field perpendicular to the line-of-sight, and therefore the spectral
index profiles shown in Fig. 15 will be shifted in log frequency
by the change in ln B⊥. Given that dβ/dln ν≈ −0.13 at around
1 GHz, the spatial uniformity of the model spectral index implies
that the average value of B⊥along each line-of-sight varies by at
most 15% (rms) from one line-of-sight to another. Of course,
A25, page 18 of 45
Planck Collaboration: Planck 2015 results. XXV.
each model spectrum is an integral of the emission along the
line-of-sight, along which B⊥varies considerably; however, the
main variation of Bis with Galactic scale height z, and hence all
lines of sight, except those at very low latitude, sample nearly
the same range of field strengths with nearly the same relative
weighting. One might expect geometric factors alone to intro-
duce significant spectral variations, but in these models the ran-
dom component of the field dominates, so to a good approxima-
tion B⊥≈√2/3|B|.
The synchrotron intensity at fixed frequency depends on B⊥
as jν∝B−(β+1)
⊥, which suggests that regions of enhanced field
strength will disproportionately dominate the emission. In the
GALPROP models, the CRLs are supposed to diffuse efficiently
through the interstellar medium, so that the local CRL density
is not correlated in detail with the field strength. In reality, if
the field is enhanced due to compression, e.g., by a shock wave,
the CRLs will respond adiabatically if the time to diffuse out of
the region is much longer than the timescale for compression.
Then the particle energies will also be shifted to higher ener-
gies, giving a larger shift in the β(ν) profile, roughly ∝B2, and
jν∝B−2β−3(the exact index depending on the field geometry, see
Leahy 1991 for details). Obviously this substantially increases
the dominance of regions of high Bin the overall synchrotron
emission.
The spatial uniformity of the modelled synchrotron spectrum
is thus likely to be an artefact due to the omission of sub-kpc
structures in the models, such as the loops and spurs that dom-
inate the high-latitude synchrotron sky. Indeed, we know that
many individual supernova remnants (SNR) have spectra sig-
nificantly flatter than that of the diffuse Galactic emission, as
expected from the steepening caused by energy-dependent diffu-
sion and radiative losses. It may seem redundant to argue on the-
oretical grounds that the GALPROP models show too little spec-
tral structure, since much larger variations have been reported
based on analysis of existing radio surveys (e.g., Lawson et al.
1987;Reich & Reich 1988;Dickinson et al. 2009). But if one
discounts features associated with free-free absorption at very
low frequencies and free-free emission at higher frequencies, the
remaining large variations are mostly in the regions of lowest
synchrotron intensity, which are susceptible to systematic errors
such as sidelobe contamination from the strong Galactic plane
emission, ground spillover, zero-level errors, and (at low fre-
quencies) pick-up of sky emission reflected from the ground.
Reich & Reich (1988) review previous large-area spectral in-
dex maps and conclude that all are unreliable, with the excep-
tion of the 38:408 MHz map of Lawson et al. (1987), which is
strikingly uniform. Reich & Reich (1988) present their own map
of spectral index between 408 and 1420 MHz in the northern
hemisphere showing extensive regions at high Galactic latitude
with β > −2.4, which they attribute at least in part to free-free
emission uncorrelated with Hα, presumably due to extinction
of the latter. They later extended their analysis to the full sky
(Reich et al. 2004), finding even more extensive flat-spectrum
regions in the southern hemisphere. If this were due to a free-
free contribution, or even to an unusually flat synchrotron power
law, the foreground temperature in the steradian or so around
(l,b)=(240◦,−40◦) would be several times higher than ob-
served at 20–30 GHz (e.g., Fig. 1a). We believe that the culprit
is likely to be some combination of the systematic errors men-
tioned above.
These considerations suggest two ways that our parameter-
ized models can allow for larger spectral variations than found
by Orlando & Strong (2013). To model emission from regions
with Bsignificantly different from the average, we can shift the
Table 3. Key parameters of the WMAP MEM and MCMC foreground
models (see Bennett et al. 2013a for details).
Model βffβsync AME
MEM −2.15 −3.0 CNM, νpeak varies
MCMC-c base −2.16 −3.0. . .
MCMC-e sdcnm −2.16 −3.0 CNM, νpeak fixed
MCMC-f fs −2.16 varies CNM, νpeak fixed
MCMC-g fss −2.16 variesaCNM, νpeak =14.95 GHz
Notes. (a)As described in Strong et al. (2011).
spectrum in ln ν, keeping the overall shape fixed. However, there
is certainly more spectral freedom allowed than this: from the
dispersion in the spectral indices of young SNR, we expect the
injected energy spectrum to be spatially and temporally variable;
moreover the propagation effects included in the synchrotron
models of Fig. 15 induce small variations in the high-frequency
spectrum, which cannot be explained by variations in B. If we
denote by TGP(ν) a fiducial spectrum from that simulation, then
we can steepen or flatten it by writing
Tsyn(ν)=(ν/ν0)δβ TGP(ν),(4)
i.e., a change of local spectral index β(ν)=dln T/dln νby a con-
stant δβ. Most of the variation in Fig. 15 can be accounted for by
such a model, but a larger range of δβ is needed to account for the
actual sky. In practice, detailed accounting for the spectral cur-
vature below 10 GHz is superfluous given that we only use one
observed frequency in that regime. For our baseline Commander
analysis our template was the spectrum of a single pixel in the
z10LMPD_SUNfE model from Orlando & Strong (2013), chosen
to be close to the all-sky median for several spectral parameters
(i.e., knee frequencies, and spectral indices at the knees). As de-
scribed in Planck Collaboration X (2016), this was fitted to the
data along with the other foreground components allowing for a
global shift in frequency (determined to be a factor of 0.26), and
an amplitude that was fitted at each pixel.
The resulting synchrotron amplitude map (shown in Planck
Collaboration X 2016) is essentially determined by the high S/N
408 MHz data, and deviates from it only in that residual free-free
emission at 408 MHz is corrected (or over-corrected, at bright
compact H ii regions in which free-free absorption is significant
at 408 MHz). The derived global frequency shift is surprisingly
large and has the effect of steepening the average spectrum be-
tween 408 MHz and the Planck bands, resulting in a relatively
low amplitude for synchrotron compared to AME and free-free.
This may be because the fit is most strongly constrained by the
North Polar spur and the diffuse halo of the inner Galaxy, the two
regions that are both strong in synchrotron and relatively free of
other components that could absorb errors in the synchrotron fit.
Most analyses find flatter spectra in the narrow Galactic plane
(e.g., Planck Collaboration Int. XXIII 2015).
4.4. Comparison with WMAP models
Here, we discuss the main similarities and differences be-
tween the low-frequency foreground component maps from our
Commander analysis of the combined Planck and WMAP data
sets (Planck Collaboration X 2016), and from the final WMAP
analysis (Bennett et al. 2013a).
The WMAP team have published sets of component-
separated maps using a maximum entropy method (MEM) and
A25, page 19 of 45
A&A 594, A25 (2016)
Table 4. Ratios of the WMAP MEM and MCMC maps to the
Commander maps for the high latitude sky with |b|>20◦(top) and for
|b|<20◦(bottom).
Run Sync Free-free AME
a r a r a r
MCMC-c base 0.50 0.62 0.68 0.77 . . . . . .
MCMC-e sdcnm 0.52 0.92 0.77 0.87 4.91 0.75
MCMC-f fs 0.52 0.62 0.80 0.77 3.18 0.67
MCMC-g fss 0.55 0.62 0.77 0.78 3.14 0.70
MEM 0.34 0.84 0.76 0.79 2.18 0.86
MCMC-c base 0.12 0.74 0.91 0.83 . . . . . .
MCMC-e sdcnm 0.61 0.99 0.67 0.98 2.71 0.91
MCMC-f fs 0.13 0.78 1.09 0.89 4.71 0.86
MCMC-g fss 0.15 0.77 1.06 0.93 3.22 0.88
MEM 0.19 0.90 1.01 0.97 2.36 0.93
Notes. A value a>1 indicates that the Commander maps contain more
emission than the WMAP maps. The Pearson correlation coefficient (r)
is also given.
Fig. 16. Scatter plot between the Commander synchrotron solution eval-
uated at 22.8 GHz and the WMAP MCMC-f synchrotron model at the
same frequency.
a Markov chain Monte Carlo (MCMC) technique on a pixel ba-
sis (Bennett et al. 2013b). We have compared the synchrotron,
free-free and AME maps produced by these methods with the
Commander maps through T–Tplots with HEALPix Nside =64.
Table 3gives the key model parameters of the WMAP models
that we used. Point sources from the PCCS2 catalogue brighter
than 1 Jy at 28.4 GHz have been masked out. The best-fit ratios
of the maps, derived from the T–Tplots, are given in Table 4
for the high-latitude sky, with |b|>20◦and for |b|<20◦;ais
the best-fitting slope. Also listed in Table 4are the Pearson cor-
relation coefficients (r) for all the comparisons. If r&0.9, the
measured slope, a, is reliable, since there is a good linear rela-
tionship between the Commander and the WMAP templates.
A first obvious difference between the Commander and
WMAP MEM/MCMC models is that the AME component is
systematically higher (between 2–4 times) in the Commander so-
lution, at the expense of the synchrotron emission that is lower in
Commander. This is clearer on the region closer to the Galactic
plane (bottom half of Table 4), due to the better correlations mea-
sured there.
When comparing the Commander synchrotron solution with
the WMAP models that allow the synchrotron spectral index to
vary over the sky, the data are not well-fitted with a single slope,
since two populations are clearly present. Figure 16 shows the
scatter plot between the Commander synchrotron solution eval-
uated at 22.8 GHz and the WMAP MCMC-f synchrotron model
at the same frequency. The two slopes visible in the figure are
the result of a flatter spectrum synchrotron component on the
plane (e.g., Kogut et al. 2007), which is not accounted by the
Commander synchrotron model, where βsync is fixed over the sky.
This is not the case for the WMAP MCMC-e “sdcnm” and MEM
solutions, which use a fixed synchrotron spectral index, resulting
in a much better correlation between the Commander and WMAP
components (see Pearson correlation coefficients for this compo-
nent in Table 4). We also note from Table 4that the Commander
synchrotron component is always lower than the WMAP syn-
chrotron models, which is due to the larger AME component in
the Commander model.
The Commander free-free component is similar to most of
the WMAP models at low Galactic latitudes, where Pearson’s
r&0.89 for all but the MCMC-c model. The correlation coeffi-
cients in this case are also close to 1. At high Galactic latitudes
(top half of Table 4), the Commander free-free solution is lower
than the WMAP models, at the expense of a higher AME con-
tribution in Commander. The wider frequency range of Planck,
specifically from 100–217 GHz, should enable a cleaner separa-
tion of the free-free component (Planck Collaboration Int. XIV
2014;Planck Collaboration Int. XXIII 2015).
In addition to the baseline Commander model, we have also
compared the WMAP models to several earlier Commander runs
with different input parameters. In particular, we look at a run
that used a broken power-law for the synchrotron component
along with a single AME component, and another that had
straight power-laws for both the synchrotron and free-free com-
ponents, also with a single AME component. In these models
the synchrotron spectral index is fitted locally, while the break,
∆β, in the first is fitted globally. This parameterisation is closer
to that used by the WMAP team than is our baseline case with
two AME components and a GALPROP synchrotron spectrum,
and the fitted results are also closer to the WMAP results, with
less AME and more synchrotron; however some differences re-
main. In both models the free-free amplitude in the Commander
outputs is still lower than in the WMAP models, and in both
Commander still finds less synchrotron emission than WMAP at
low latitudes (except for WMAP “sdcnm”). However, at high lat-
itudes they find more power in the synchrotron component than
does the WMAP analysis. The AME components in these ear-
lier Commander models are weaker than in our baseline and in
much closer agreement with WMAP, with ratios around one for
both the “fs” and “fss” models at both low and high latitudes,
while the Commander solutions find more high-latitude and less
low-latitude emission than the “sdcnm” model. These runs were
not used for the baseline product as they gave worse overall fits
and are less physically motivated than our eventual choice; we
note them here solely to qualitatively consider the effects that
different models can have on the component outputs.
The WMAP model that is closest to the baseline Commander
solution is MCMC-e “sdcnm”, although the AME component
is brighter in the Commander products. Also, the Commander
free-free amplitude is about 30% fainter than the correspond-
ing WMAP fit; this has been noted before from comparison with
RRL data by Alves et al. (2010) and Planck Collaboration Int.
XXIII (2015). The main limitation in obtaining accurate com-
ponents maps is the precise quantification of the synchrotron
A25, page 20 of 45
Planck Collaboration: Planck 2015 results. XXV.
component. Fuskeland et al. (2014) and Vidal et al. (2015) have
shown that the polarized spectral index of synchrotron emis-
sion varies across the sky and these variations should be taken
into account to obtain an accurate separation between the syn-
chrotron and AME components; low frequency (5–20 GHz) data
will be crucial for this.
4.5. LMC and SMC
The Large and Small Magellanic Clouds (LMC and SMC)
are satellite galaxies of the Milky Way, close enough
(50.0±1.1) kpc for the LMC, Pietrzy ´
nski et al. 2013, and
(61 ±3) kpc for the SMC, Hilditch et al. 2005) that they are
resolved by Planck both at 50in thermal dust and 300at low
frequency. They have already been well studied at Planck and
WMAP frequencies (Bot et al. 2010;Planck Collaboration XVII
2011;Draine & Hensley 2012). They thus provide a good test-
bed for comparing the Commander solution with previous results
and expectations, both in the Magellanic clouds and in compar-
ison with our own Galaxy. We show maps of the Commander
solution in the LMC in Fig. 17, which we discuss component
by component below. The brightest region in all of the non-
CMB component maps is the Tarantula nebula, also known as
30 Doradus, a well-known H ii and star-formation region located
at (l,b)=(279.
◦5,−31.
◦7).
To cross-check the Commander solution, we perform aper-
ture photometry on both the Commander maps and the lat-
est Planck and WMAP maps, as well as other ancillary
data, to generate spectral energy distributions (SEDs) of the
LMC and the SMC. For the LMC we use central coordi-
nates of (l,b)=(279.
◦5,−33.
◦5), with a circular aperture of ra-
dius 3000, and an outer annulus for background subtraction
of 300–3500. For the SMC we use central coordinates of
(l,b)=(302.
◦8,−44.
◦3), with a circular aperture of radius 1500,
and an outer annulus of 150–2000. We do not attempt to remove
Galactic foregrounds from the maps or SEDs, except through the
background annulus in the photometry, since we are using the
SEDs for the purpose of comparison. As a result, some Galactic
foreground emission can be seen in the resulting LMC thermal
dust map in Fig. 17.
We use the same aperture photometry code that
was developed for Planck Collaboration XX (2011) and
Planck Collaboration Int. XV (2014). We have modified the
code to run on the Commander component maps to measure the
amplitudes of the individual components in the aperture. The
uncertainties are derived from the background annulus only:
they do not include model or calibration uncertainties. Results
are given in Table 5. We then sum those model SED components
to compare with the flux densities measured directly from the
Haslam, Planck, and WMAP maps, along with radio maps at
1.4 GHz by Reich et al. (2001) and 2.3 GHz by Jonas et al.
(1998) and the infrared maps from COBE-DIRBE (Hauser et al.
1998) and IRAS (Miville-Deschênes & Lagache 2005). This
comparison for the LMC and SMC is shown in the left panel of
Fig. 18. We note that the Commander solution does not include
data points above 857 GHz in its fit. In deriving the models, we
exclude the 100 and 217 GHz Planck maps, since those bands
contain CO emission.
Using the same spectral component model described in
Planck Collaboration Int. XV (2014), including colour correc-
tions, we also perform least-squares fits (LSF) to the flux den-
sities up to 3 THz. There are two variants of the LSF model re-
ported in Table 5and Fig. 18. One is fitted to data including
the CMB contribution (henceforth referred to as “LSF”) and the
Table 5. Values of the fitted parameters for the LMC (top) and the SMC
(bottom) from Commander, and the least-squares fitting before and after
CMB subtraction.
Parameter Commander LSF (CMB) LSF (no CMB)
EM 66.2±1.7 57 ±4 67 ±2
Te[K] 7000 8000 8000
Ad24 903 ±488 (3.1±0.6) ×10−5(2.82 ±0.14) ×10−5
βd1.516 ±0.003 1.47 ±0.07 1.42 ±0.03
Td[K] 19.29 ±0.09 22.7±0.9 23.2±0.3
AAME1 15.3±0.3. . . . . .
AAME2 3.47 ±0.14 . . . . . .
ACMB [µK] 16 ±3 15 ±3. . .
Async [Jy] 239 ±4 406 ±33 375 ±31
βsync −3.1−2.70 ±0.05 −2.74 ±0.04
EM 20.0±1.8 21.7±0.9 25.1±0.4
Te[K] 7000 8000 8000
Ad2541 ±70 (1.26 ±0.07) ×10−5(7.2±0.3) ×10−6
βd1.463 ±0.004 1.36 ±0.05 1.06 ±0.02
Td[K] 18.57 ±0.14 21.5±0.3 25.35 ±0.18
AAME1 1.497 ±0.07 . . . . . .
AAME2 0.57 ±0.03 . . . . . .
ACMB [µK] 36 ±5 37 ±2. . .
Async [Jy] 25 ±2 34 ±3 31 ±3
βsync −3.1−2.76 ±0.11 −2.88 ±0.11
Notes. The Commander parameters are as described in Planck
Collaboration X (2016) except for Async, which has been rescaled to a
reference frequency of 1 GHz. We note that the Commander uncertain-
ties only include the standard deviation, and do not include modelling
uncertainties. The LSF parameters are the same, except for Ad, which is
the optical depth at 250 µm.
other is fitted to data post-CMB subtraction, namely flux densi-
ties from aperture photometry on CMB-subtracted maps (using
the Commander CMB solution; henceforth referred to as “LSF-
CMB”).
We now go through each component in turn for both the
LMC and SMC, considering both the morphology and the SEDs.
CMB. The CMB maps for both the LMC and SMC do not ap-
pear to be strongly contaminated by foreground emission; there
is no clear correlation between the CMB map and the fore-
ground components. The average CMB contribution agrees well
in both the Commander and LSF estimates in the LMC and SMC.
CMB contributions in both objects are positive, as was seen in
Planck Collaboration XVII (2011). We thus conclude that the
Commander CMB solution is robust at separating the CMB from
foreground emission in this region.
Synchrotron. Low-frequency synchrotron emission is present
in both the LMC and SMC. It is subdominant in total intensity
at Planck and WMAP frequencies. In the LMC the peak of the
emission is slightly offset from the Tarantula nebula, and dif-
fuse emission is also present. The Commander spectral index be-
tween 408 MHz and WMAP/Planck frequencies is assumed to
be βsync =−3.1 according to the GALPROP model (see Sect. 4.3);
in the LSF and LSF-CMB fits for the LMC we find βsync ≈−2.7,
and for the SMC βsync ≈ −2.8, effectively between 408 MHz and
22.8 GHz. At similar frequencies, Israel et al. (2010) found that
βsync =−2.70 ±0.05 in the LMC; however, they found a steeper
index of −3.09 ±0.10 in the SMC that agrees better with the
Commander model. As such, Commander will under-estimate the
A25, page 21 of 45
A&A 594, A25 (2016)
(a) Synchrotron
10 20 30 40 50
µK at 30 GHz
(b) Free-free
0 10 100 1000
cm−6pc
(c) Spinning dust
10 30 100 300
µK at 30 GHz
(d) Thermal dust
10 100 1000
µK at 545 GHz
(e) Synchrotron polarization
0 10 20 30
µK at 30 GHz
(f) CMB
−200 −100 0 100 200
µKCMB
(g) CO J=2→1
0 0.2 0.4 0.6 0.8 1.0
K km s−1
(h) Thermal dust polarization
0 5 10 15 20
µK at 353 GHz
Fig. 17. Commander solution in the Large Magellanic Cloud region, plotted in Galactic coordinates. Panels show from left to right and top to
bottom:a) synchrotron brightness temperature at 30 GHz with 1◦resolution (linear colour scale); b) free-free emission measure with 1◦resolution
(logarithmic colour scale); c) spinning dust brightness temperature at 30 GHz with 1◦resolution (logarithmic colour scale); d) thermal dust
brightness temperature at 545 GHz with 50resolution (logarithmic colour scale); e) synchrotron polarization amplitude, P, at 28.4 GHz with 1◦
resolution (corrected for polarization leakage, linear colour scale); f) CMB temperature with 50resolution (linear colour scale); g) CO J=2→1
emission with 50resolution (linear colour scale); and h) thermal dust polarization amplitude, P, at 353 GHz with 50resolution (corrected for
polarization leakage, linear colour scale). Each map covers 15◦×15◦, and is centred on Galactic coordinates (l,b)=(279◦,−34◦).
101
102
103
104
105
106 0.1 1 10 100 1000
Flux density (Jy)
Wavelength (mm)
Model
Archival data
WMAP
Planck
101
102
103
104
105
106 0.1 1 10 100 1000
Flux density (Jy)
Wavelength (mm)
Model
Archival data
WMAP
Planck
101
102
103
104
105
106 0.1 1 10 100 1000
Flux density (Jy)
Wavelength (mm)
Model
Archival data
WMAP
Planck
100
101
102
103
104
105
0.1 1 10 100 1000
Flux density (Jy)
Frequency (GHz)
Model
Archival data
WMAP
Planck
100
101
102
103
104
105
0.1 1 10 100 1000
Flux density (Jy)
Frequency (GHz)
Model
Archival data
WMAP
Planck
100
101
102
103
104
105
0.1 1 10 100 1000
Flux density (Jy)
Frequency (GHz)
Model
Archival data
WMAP
Planck
Fig. 18. SEDs of the LMC (top) and SMC (bottom) using aperture photometry, from the Commander solution (left), and least-squares fits to
the data with (middle, LSF) and without (right)) the CMB (LSF-CMB). The components are synchrotron (orange dashed lines), free-free (black
double-dashed lines), AME (green dashed lines), CMB (blue dashed lines), and thermal dust (red dashed lines). The solid black line represents the
sum of the model components. Planck data points are shown in red, WMAP in blue, and other ancillary data (some of which are not included in
the fits, see text) in black.
A25, page 22 of 45
Planck Collaboration: Planck 2015 results. XXV.
synchrotron contribution at Planck frequencies, particularly in
the LMC and to a lesser extent in the SMC. The synchrotron
amplitudes at 1 GHz in Commander for both the LMC and SMC
are significantly lower than that in the LSF, due to the steep
spectrum assumed, with the difference absorbed by the free-free
component. The LMC also appears in the synchrotron polariza-
tion map; the polarized emission is at its highest (around 30%
polarized, although this would be lower if the Commander syn-
chrotron intensity spectral index were flatter) to the left of the
peak in the synchrotron total intensity, offset from Tarantula,
which corresponds to the two polarized synchrotron filaments
identified by Klein et al. (1993).
Free-free emission. The free-free component dominates the
SEDs of both the LMC and the SMC at frequencies of 5–
50 GHz. The majority of the emission comes from the Tarantula
nebula; there is also diffuse emission closer to the centre of the
LMC, and other compact (<1◦) objects are present. The mor-
phology agrees well with the Hαmaps of individual sources
in the LMC by Davies et al. (1976, DEM): the main region
of 30 Dor is surrounded by a large number of smaller sources
that show up in the Commander map as diffuse emission. The
two lower regions are groups of sources, with one compris-
ing DEM 4, 6, and 36 (bottom-left), and the other consisting
of DEM 27, 28, and 29 (bottom-right). The amplitude of the
free-free emission can be converted to a star-formation rate
(SFR) using the equations in Condon (1992). The Commander
free-free amplitude at 10 GHz is 141 Jy, which gives an SFR
of 0.10 Myr−1; the LSF model gives 0.08 Myr−1, and the
LSF-CMB model gives 0.09 Myr−1. These are lower than
the estimate of 0.2 Myr−1for the average SFR from anal-
yses of stellar populations by Harris & Zaritsky (2009) and
Rezaeikh et al. (2014), but agree well with the recent star for-
mation rate of 0.06 Myr−1calculated by Whitney et al. (2008)
based on young stellar objects; these authors also give a range of
SFR estimates of 0.05–0.25 Myr−1from infrared and Hαdata.
However, all of these estimates correspond to an SFR over dif-
ferent timescales (e.g., Murphy et al. 2012), as well as being sub-
ject to systematic effects; a more detailed study would be neces-
sary to disentangle these effects. For the SMC, the amplitude at
10 GHz of 11 Jy from the Commander and LSF models gives an
SFR of 1.1×10−2Myr−1; the LSF-CMB model gives a slightly
higher value of 1.3×10−2Myr−1from a flux density of 12.8 Jy.
These estimates of the SMC free-free amplitude agree well with
those from Draine & Hensley (2012), who attributed 11 Jy to the
free-free component at 10 GHz, in very good agreement with the
Commander amplitude.
AME. The Commander solution finds a small component of
AME in both the LMC and SMC. In the LMC, the AME
map has a bright region centred on 30 Dor, where there are
many sources seen in the higher-resolution Hαand radio
maps (McGee & Newton 1972,Davies et al. 1976), and where
the brightest emission is seen in the LMC in thermal dust.
The AME towards 30 Dor has a higher peak frequency in the
Commander solution than that on either side, following the pat-
tern of H ii regions having higher peak frequencies, as noted
in Sects. 3and 4.2. At 20 GHz, the peak of the AME com-
ponent, AME contributes about 16 Jy to the model, compared
with the free-free amplitude of around 130 Jy. In the SMC,
1.6 Jy is attributed to AME, compared with 10 Jy for the free-
free component. In both cases, AME is about 10% of the
free-free amplitude, and it is comparable to the component
separation uncertainty in the free-free component, and the ampli-
tudes of the synchrotron and CMB components at that frequency.
The ratio of free-free to AME here is significantly higher
than that seen in the Galaxy (Planck Collaboration XX 2011;
Planck Collaboration Int. XXIII 2015) due to the presence of
30 Dor, which dominates the free-free emission in the LMC. The
LSF analysis includes an AME model, but is consistent with lit-
tle or no AME (the AME contribution seen in the LSF SMC SED
has a significance below 1σ). Nor was AME seen in the LMC
and SMC SEDs in Planck Collaboration XVII (2011). Draine &
Hensley (2012) included an AME component in their SED fits to
the SMC of 3–5 Jy at 40 GHz; the Commander solution returns
much lower AME amplitudes in the SMC of 0.2 Jy at 40 GHz.
We calculate AME emissivities for the LMC for a 15◦×15◦re-
gion centred on (l,b)=(279◦,−34◦), and for the SMC in a 5◦×5◦
region centred on (l,b)=(302.
◦8,−44.
◦3). The emissivities are
given in Table 2. For the LMC we find that the AME/545 GHz
and AME/τ353 emissivities are comparable to those from the
AME regions; however, the AME/100 µm emissivity is substan-
tially lower, which is likely because the LMC has a higher dust
temperature than the Galactic average. For the SMC, we find
much lower values for all three emissivities. The LMC has com-
parable PAH levels to our Galaxy (Bernard et al. 2008), while the
SMC has lower levels, which could explain why we find higher
AME emissivities in the LMC than the SMC. The SMC has a
very small grain (VSG) population that can be seen particularly
clearly at 70 µm (Bernard et al. 2008); given the emissivities here
this VSG emission is unlikely to be connected to AME, as AME
is typically thought to be due to PAHs. We thus conclude that
the small amount of AME found by Commander in the LMC is
at the level that would be expected from our Galaxy, although
both the LMC and SMC regions may be contaminated by leak-
age from the free-free and synchrotron components, as well as
being potentially influenced by the steep synchrotron spectral in-
dex assumed in the Commander model. Further work is needed to
improve constraints on the AME and its properties in this region.
CO. The Commander CO J=1→0 map detects various regions
of CO emission, particularly around the 30 Dor region, as well as
in various other dusty regions in the LMC. The morphology of
the map compares well with ground-based surveys, such as the
Magellanic Mopra Assessment (Wong et al. 2011). This indi-
cates that the separation of CO emission has worked well in this
region. An exception is a large negative region in the Commander
high-resolution map, which is a component separation artefact.
Thermal dust. The thermal dust seen by Planck agrees well
in morphology with IRAS data. The dust temperatures from
Commander are significantly lower than those from the LSF:
for the LMC the average dust temperature in the aperture is
(19.29 ±0.01) K, compared with (22.7±0.9) K for the LSF.
Planck Collaboration XX (2011) found (21.0±1.9) K when fit-
ting for a free βd=1.48 ±0.25 (comparable with the βdfound
here). Commander only fits data up to 857 GHz and the fitted
temperature and spectral index underestimates the flux densities
at higher frequencies. The LSF fits data up to 3 THz and so finds
a higher temperature. We also note that where the fitted tem-
perature is higher, βdis lower. The same applies for the SMC,
where the average from Commander is (18.57 ±0.14) K, com-
pared to (21.5±0.3) K from the LSF, and (22.3±2.3) K from
Planck Collaboration XX (2011), where βd=1.21 ±0.27 is flat-
ter than found by Commander but in better agreement with the
A25, page 23 of 45
A&A 594, A25 (2016)
LSF9. In the SMC, the submm excess (see e.g., Israel et al. 2010,
Planck Collaboration XVII 2011) has been subsumed into the
dust spectral index and the lower dust temperature. The thermal
dust polarization in the LMC traces the spiral structure, with the
projected magnetic field running parallel to the structure, partic-
ularly in the 30 Dor region where the Planck data have a higher
sensitivity due to its scan strategy. This is also where the syn-
chrotron polarization is seen.
In conclusion, we find that the Commander component maps in
the regions of the LMC and SMC largely agree well with previ-
ous results and expectations, although the thermal dust proper-
ties are only representative up to 857 GHz.
5. Polarized foregrounds at λ≈1 cm
The previous section showed that significant uncertainties re-
main for component separation in total intensity. Fortunately, the
picture is quite different in polarization, where, between Planck
and WMAP, we have twelve bands with maps of the all-sky po-
larization. Our current understanding is that there are three sig-
nificantly polarized components: the CMB; synchrotron emis-
sion; and thermal dust. AME polarization has not been clearly
detected, with upper limits of typically a few per cent (see
Sect. 5.7).
The three strongly polarized components have radically dif-
ferent spectra, and moreover, CMB polarization is well separated
in angular scale from the foregrounds: after convolution to 1◦
resolution, the CMB polarization predicted by the Planck best-
fit cosmology contributes an rms of only 0.54 µK to the Qand U
maps, negligible compared to the noise per beam (although, of
course, detectable in the angular power spectrum at low multi-
poles). The spectrum of the polarized dust is discussed in Planck
Collaboration Int. XXII (2015) and Planck Collaboration Int.
XXX (2016). Here we are interested primarily in the synchrotron
component and consider the dust polarization only to the ex-
tent that the magnetic field pattern that it traces sheds light on
the synchrotron features (Sect. 5.2). Planck Collaboration Int.
XIX (2015) informs our understanding of the dust polarization
features.
We begin by constructing a new map of polarized syn-
chrotron emission by combining WMAP and Planck maps,
which significantly increases the average S/N compared with
WMAP/Planck data alone. We then discuss some of the major
features in the polarized sky, including the loops, spurs, fila-
ments, and bubbles.
5.1. Combination of Planck and WMAP
WMAP (Page et al. 2007;Bennett et al. 2013a) provided the
first clear view of the intrinsic synchrotron polarization across
the sky – previous ground-based observations (e.g., Brouw &
Spoelstra 1976;Wolleben et al. 2006) being strongly affected by
Faraday rotation and depolarization due to their lower observ-
ing frequencies (.2 GHz). WMAP’s lowest-frequency channel,
K-band, dominates the combined WMAP S/N in polarization be-
cause of the steep spectrum of synchrotron emission. Although
9The Commander solution has a Gaussian prior on the thermal dust
spectral index of β=1.55 ±0.1, and the fitted values for the LMC
and SMC are towards the lower end of this prior; however Commander
would have returned a lower value for βthan the prior if the data had
preferred it. The LSF indicate flatter spectral indices, however this is
including data over a larger (>857 GHz) range of frequencies.
Fig. 19. Difference between the combined weighted polariza-
tion maps for Planck and WMAP at 1◦resolution, defined as
h(QPlanck −QWMAP)2+(UPlanck −UWMAP )2i1/2.
the full-mission Planck-LFI maps have significantly lower noise
than the final WMAP results, the foreground S/N is very similar,
since at Planck’s lowest frequency (28.4 GHz) the synchrotron
brightness is half that in K-band. After smoothing to a com-
mon resolution of 1◦, the median (mean) S/N is 2.47 (3.77) for
WMAP K-band and 2.64 (3.72) for Planck 28.4 GHz. The differ-
ent scan strategies of the two missions result in somewhat differ-
ent sensitivity patterns, so that each map is superior in some sky
regions; Planck has the largest advantage in the regions around
the Ecliptic poles (see the hit count maps in Planck Collaboration
VI 2016).
Planck and WMAP polarization results generally agree well,
but the difference map (see Fig. 19) shows significant artefacts,
both on the plane and at high latitudes. Extremely large angular-
scale (`=3,5,7) residuals covering much of the sky are thought
to be due to the poorly-constrained modes in the WMAP data as
a result of the scan strategy (Jarosik et al. 2011;Bennett et al.
2013a); they occur in all the WMAP frequency maps, with an
approximately consistent pattern. However, they are most readily
visible in the V- and W-bands of WMAP when the foreground
brightness is close to the minimum. The on-plane features are
likely to be caused by inaccuracies in the Planck leakage maps.
These differences are further discussed in Planck Collaboration
X(2016).
Because polarization is only marginally detected in many
pixels of the individual WMAP and Planck maps, we have con-
structed a weighted mean image using all the maps in the range
20–50 GHz, where the polarization is dominated by optically
thin synchrotron radiation, with negligible Faraday rotation (ex-
cept at the Galactic centre, Vidal et al. 2015), so that the polar-
ization angles should be consistent between the maps. We first
smooth the Planck and WMAP maps to a resolution of 1◦and
Nside =256 in both Qand U, and the QQ,UU, and QU covari-
ances as well. We assume that the emission has a single spectral
index of −3.0, and we use this to scale all the maps (and their
covariances) to match the Planck 28.4GHz image. We then use
the 2 ×2Qand Ucovariance matrices and their inverse ma-
trix to calculate the weighted average in each pixel of the Q
and Umaps. We create three weighted maps as a result: one
consisting of the three lowest WMAP bands; one consisting of
the two lowest Planck bands; and one using both WMAP and
Planck. The Planck variance estimates include the uncertainty
in the bandpass leakage correction. This results in a large frac-
tional error in Qand Ualong the plane, so WMAP data dominate
the on-plane emission in the combined map. At high latitudes,
A25, page 24 of 45
Planck Collaboration: Planck 2015 results. XXV.
the somewhat higher sensitivity of Planck down-weights the im-
pact of the large-scale WMAP polarization artefacts, but we note
in our discussion when these may still have an impact on our
results. The improvement in S/N is significant, with: a median
(mean) S/N of the WMAP weighted map of 2.70 (4.19); the
Planck weighted map of 2.68 (3.81); and the combination of
both of 3.37 (5.39), where the bigger improvement comes from
the combination of both WMAP and Planck. This corresponds
to a roughly 25% improvement in S/N compared to using the
combined WMAP or Planck maps separately.
Commander also provides a polarized intensity synchrotron
map, which is calculated using only Planck data. As this map is
the best-fit solution to a model using noisy data, it is noisier than
our combined Planck-only map (our map treats the small CMB
component as noise). Moreover, the addition of WMAP K-band
data in our map further reduces the noise in the final combined
map. At 1◦resolution, the Commander map has a median noise
value over the entire sky of σQ≈σU=4.5 µK, while our com-
bined map has σQ≈σU=2.6 µK.
Polarized intensity maps have been debiased using the
asymptotic estimator (Montier et al. 2015;Vidal et al. 2016),
which generalizes the estimator first proposed by Wardle &
Kronberg (1974) to the case of anisotropic errors in (Q,U).
Figure 20 shows the resulting polarization combination maps. A
number of new polarized structures visible in the combined map
that were unclear in the individual maps are highlighted. These
are discussed in the sections that follow.
5.2. Overview of polarized synchrotron emission
Synchrotron total intensity is distributed comparatively uni-
formly over the sky. In the 408 MHz map, assuming an instru-
mental plus extragalactic background of 8.9 K (Wehus et al.
2016), half the total Galactic flux is contributed by 18% of
the sky10. The equivalent figure for the dust-dominated Planck
545-GHz map is 4.6%. In the synchrotron intensity map, al-
though the Galactic plane and the North Polar spur are visually
prominent, they are superposed on broad diffuse emission that
dominates the total flux.
This diffuse emission seems to be much weaker in polar-
ization. Away from the narrow Galactic plane, i.e., at |b|&3◦,
the polarized cm-wavelength emission is dominated by the syn-
chrotron loops and spurs familiar from low-frequency radio sur-
veys. In fact, Vidal et al. (2015) demonstrate a close correspon-
dence between the polarized intensity at WMAP K-band and an
unsharp-masked version of the 408 MHz map, in which structure
on scales &10◦is filtered out (see Fig. 21 top row a and b).
In particular, there is hardly any trace in the Planck and
WMAP polarized maps of the synchrotron halo of the inner
Galaxy, which fills roughly |l|.60◦and 5◦<
∼|b|<
∼15◦
(Planck Collaboration Int. XXIII 2015) and contributes 20–30%
of the Galactic synchrotron flux in low-frequency maps (com-
pare panels a and c in Fig. 21). The effect is even clearer in
fractional polarization, which is very low (.10%) in the inner
halo region (see Fig. 22). The component separation analysis
of Sect. 4implies that the halo is still present in the unpolar-
ized cm-wavelength maps, because the data are well-fitted with
our synchrotron model, which has a constant shape, so that the
synchrotron distribution is no different at 20 or 30 GHz from
10 The diffuse emission may be even brighter, since the background as-
sumed may include an isotropic component of Galactic emission. Using
the Lawson et al. (1987) extragalactic background estimate of 5.9 K,
half the sky flux comes from 21% of the sky.
408 MHz. Although, as we have seen, this separation is subject
to substantial uncertainties, it is corroborated by spectral analy-
sis of ground-based surveys (Reich & Reich 1988;Platania et al.
2003) that find that the halo and spurs have similar spectral in-
dices near 1 GHz.
The magnetic field in the Galactic disc follows the plane, as
expected from the shearing effect of differential rotation, and one
would expect the same to apply in the inner halo. Its polariza-
tion would then be roughly orthogonal to that in the spurs rising
out of the plane, which all have fields roughly parallel to their
axes (Vidal et al. 2015). Therefore, the net polarization of halo
and spurs will cancel to some extent. But the fact that the spurs
are prominent in polarization, even when they are only perturba-
tions on the inner halo in total intensity (see Fig 23), implies that
the halo must be very weakly polarized; even between the spurs,
there is hardly any sign of a field parallel to the plane in the inner
halo region (a trace may be visible near (l,b)=(34◦,8◦); see the
discussion of the l=45◦feature in Sect. 5.4). This is a surprising
contrast to the narrow plane at |b|.3◦, where the overall par-
allel orientation of the field is very clear in Fig. 20, despite the
fact that a significant fraction of its emission is due to individ-
ual SNRs (Planck Collaboration Int. XXIII 2015), whose over-
all polarization orientation is nearly random. The conventional
analogy between the Galactic halo and the Solar corona suggests
that the halo field should relax to a nearly force-free configura-
tion with little small-scale structure, constrained primarily by the
foot points where the field lines are tied to dense gas clouds in
the plane. This would lead us to expect a reasonably high frac-
tional polarization; differential rotation of the footpoints should,
as with the disc field, shear the field so that it is largely parallel to
the plane as viewed from Earth. Evidently the halo field is much
more tangled than this naive argument would suggest.
It is worth noting that the overall fractional polarization also
appears to be low (<15%) across the high latitude sky (Fig. 22),
although this quantity is very sensitive to the poorly-known zero
level of the synchrotron total intensity (Vidal et al. 2015). Again,
the implication is that there is substantial tangling of the field
even on lines of sight looking out of the disc.
5.3. Loop I
5.3.1. Structure
Loop I is the nearly-circular structure of radius 58◦whose top-
left quadrant (as viewed in Galactic coordinates)11 is traced by
the North Polar spur (NPS). It is also detected in soft X-rays,
where the emission is dominated by thermal emission from
3×106K gas (Willingale et al. 2003), and is bordered by cold ma-
terial visible via H iand dust emission (both thermal and anoma-
lous). The NPS has long been a suspected γ-ray emitter, and
Ackermann et al. (2014) clearly detect Loop I at GeV energies,
presumably due to inverse-Compton scattering of starlight by the
CRLs in the loop, combined with pion decay emission from the
cold border. Figure 21 shows the structure in several tracers in
stereographic projection (chosen because circles are projected
as circles and also because the angular scale increases outwards,
enabling a clear display of features around the edge of this enor-
mous structure that covers most of a hemisphere).
In both Galactic hemispheres the polarization maps show a
number of spurs within Loop I that parallel its outer boundary,
11 To avoid the non-intuitive concepts of Galactic “East” and “West”,
throughout this section we describe emission features in terms of left
and right as projected in the figures, corresponding to the directions of
increasing and decreasing longitude, respectively.
A25, page 25 of 45
A&A 594, A25 (2016)
40
0
Fig. 20. Top: combined weighted polarization intensity map after debiasing, with features highlighted. The black dash-dot lines show the outlines
of Loops I to IV, as defined by Berkhuijsen et al. (1971), the blue dashed lines show the filaments described by Vidal et al. (2015) using WMAP
polarization data, the red dashed lines show features that are visible in the new Planck data, and the magenta dashed lines show the outline of
the Fermi bubbles. Filaments that are discussed in the text are labelled. Bottom: the same combined polarization intensity map, with projected
magnetic field angle (at 90◦to the polarization angle) encoded in colour with asinh scaling. The coloured half-disc represents the polarization
angle depicted in the map, with 0◦horizontal, while the polarization intensity is represented with the radial distance along the half-disc.
nearly all concentrated on the left-hand side of the structure.
Since Loop I covers about a third of the sky, including the inner
Galaxy, some of the features projected inside it are surely unre-
lated, but the general coherence of the structure strongly suggests
that most of the emission away from the plane has a common
cause. The internal spurs are much more obvious in polarization
than in total intensity, even after unsharp masking (Figs. 21a–
c, 23). These features are therefore highly polarized; however,
they are projected onto diffuse extended emission that seems
to be weakly polarized (.15%), including the inner halo of the
Galaxy, so the overall fractional polarization towards the inner
spurs is not much higher than towards the NPS itself (≈40–50%),
while between them the fractional polarization is low (Fig. 22).
Previous measurements of the radio outline found Loop I
to be surprisingly close to circular: the definitive study by
Berkhuijsen et al. (1971, hereafter BHS) found that 19 points
on the ridge-line covering 155◦around the loop (all in the north
Galactic hemisphere) fitted a small circle with an rms of only
0◦
.9. The NPS conforms roughly to the brightness profile ex-
pected from a shell of emission, for which the ridge-line marks
the tangent to the inner surface; the outer boundary lies several
degrees beyond the ridge line, and is also clearly traceable, espe-
cially around the NPS (e.g., Fig. 21b). Relatively bright diffuse
polarized emission outside the boundary of the NPS south of
b=40◦(Fig. 21c, 25) seems to be part of a different structure
(see Sect. 5.4).
A25, page 26 of 45
Planck Collaboration: Planck 2015 results. XXV.
Fig. 21. Various tracers of the interstellar medium in the hemisphere containing Loop I, with all maps centred at l=330◦,b=0◦, and using
stereographic projection, so that small circles in the sky are projected as circles: a) 408 MHz; b) unsharp-mask 408 MHz; c) combined Planck-
WMAP polarization intensity; d) Hiat vLSR =0, velocity width 10 km s−1, from the LAB survey (Kalberla et al. 2005); e) the same, at −10 km s−1;
f) Planck 353 GHz temperature; g) Planck 353 GHz polarization intensity; h) ROSAT bands R4+R5 (0.44–1.21 keV) from Snowden et al. (1997);
and i) Fermi 1–10GeV from Ackermann et al. (2014). The dashed white outline is our proposed outer boundary, derived from panels b) and c).
A25, page 27 of 45
A&A 594, A25 (2016)
Fig. 22. Polarization percentage of the debiased weighted polarization
intensity map against the Commander synchrotron total intensity fit eval-
uated at 22.8 GHz. Uncertainties in the component separation discussed
in the text, and also uncertainty in the absolute zero level, make this
highly uncertain, except in regions of strong synchrotron emission, such
as the North Polar spur and the inner Galactic halo (see also discussion
in Vidal et al. 2015).
South of the plane, the loop is superimposed on the inner
halo and cannot be clearly followed in total intensity, but in
WMAP and Planck polarization maps (Figs. 20,21), the pat-
tern is clear, at least on the left-hand side of the structure. The
brightest southern spur, called filament Is by Vidal et al. (2015),
runs well inside the path of the BHS small circle. However, a
fainter spur extends from (l,b)=(22◦,−5◦), closely along the
BHS circle until its ridge becomes lost in faint diffuse emis-
sion near b=−37◦; at this point clearly-detectable polariza-
tion again extends several degrees outside the circle. There is
then no trace of a ridge-line close to the BHS circle along its
whole bottom-right quadrant. It eventually re-aligns with the
observed loop boundary at the top of the spur north of the
Vela SNR (l,b)=(268◦,25◦), near the first measured point
in Berkhuijsen et al. (1971). In this bottom-right quadrant three
broad, faint, concentric arcs extend from the interior of Loop I,
reaching some 30◦beyond the BHS circle; these arcs are vis-
ible both in polarization and in the unsharp-masked 408 MHz
map. The outermost (filament XII of Vidal et al. 2015) rejoins
the plane near l=250◦and appears to emerge on the other side
as the aforementioned spur above Vela. Although they might be
unrelated, they share the general pattern of the other Loop I fil-
aments, for instance the middle arc that passes in front of the
LMC is clearly brightest on the left-hand side, as it approaches
the Galactic plane near l=300◦. We consider it plausible that
filament XII represents the outer boundary of Loop I, and on
this basis we have marked the outline of the loop in Fig. 21.
None of the current theories (see Sect. 5.3.3) predict a strictly
spherical structure; taking the smallest plausible distance to the
centre as 120 pc, the diameter is at least twice the 100 pc scale
height of the H idistribution. Therefore, even ignoring the pos-
sible interaction with the Local Cavity (LC) in which the Sun
is situated, the loop is expanding into an inhomogeneous envi-
ronment. Most individual SNRs, even though much smaller than
Loop I, show quite large departures from circular outlines, often
exceeding the level of distortion implied if filament XII traces
the edge of Loop I.
In X-rays, the North Polar Spur peaks at a position inwards
from the synchrotron ridge, but, as the outline drawn in Fig. 21
shows, the outer edge of the spur in the two tracers is coincident.
This result is unexpected, because the Spur cannot be bounded
by a shock front fast enough (v&300 km s−1) to heat the ambient
Fig. 23. Profiles of Loop I in total intensity at 408 MHz (Remazeilles
et al. 2015) and in our combined Planck-WMAP polarized intensity im-
age at 28.4 GHz. Profiles are along radial cuts beginning at the nominal
centre of the loop, (l,b)=(329◦,17◦
.5) and are averaged over a 10◦-wide
sector in position angle, at position angles 75◦(solid line), 50◦(dotted),
and 25◦(dashed). They cross the North Polar spur at b=22◦,43◦, and
61◦, respectively.
medium to millions of kelvin, while at the same time accelerat-
ing the cold border to at most 25 km s−1(see Sect. 5.3.3 below).
In young supernova remnants, co-incidence of radio and X-ray
boundaries is always due to X-ray synchrotron emission from
the shock, and it seems likely that in the NPS the faint X-ray
emission from the outer edge is also non-thermal.
In contrast, the gas and dust tracers extend a few degrees
further out, particularly obvious in the H iand dust filament at
the top of the loop, (l,b)=(320◦,84◦). This is just as expected,
since neutral atoms and dust grains could not survive at the 3 ×
106K of the X-ray emitting gas.
5.3.2. Distance
Loop I is usually associated with the Sco-Cen OB association
(e.g., Salter 1983) at a distance D=120–140 pc (de Zeeuw et al.
1999). The primary evidence is that the NPS appears to be de-
tectable in starlight polarization for stars at distances of &100 pc
(Mathewson & Ford 1970;Santos et al. 2011); the aligned grains
must be in the cold border rather than the spur itself. Planck-HFI
maps of polarized dust (Fig. 21g and Planck Collaboration Int.
XIX 2015) show that emission and extinction measurements of
field direction agree, confirming that the extinction distances ap-
ply to the dust features seen in emission.
Sofue (1977,1994) has instead suggested that the loop is due
to an outburst at the Galactic centre, but in Sect. 5.5 we argue
that this model is now ruled out. However, evidence against the
traditional distance has also been accumulating.
Iwan (1980) showed that the wall of dense gas that bounds
the LC would have to form a major “dent” in the apparently
spherical structure of Loop I, assuming D=130 pc. Placing the
loop several hundred parsecs away avoids the awkwardness that
this dent is not reflected at all in the projected outline. If the
shape is quasi-spherical, its near surface would only be about
20% of the distance to the centre, so D≈400 pc given absorp-
tion at the LC wall. It would still be plausible that stars at about
100 pc could be polarized by dust in the border.
A25, page 28 of 45
Planck Collaboration: Planck 2015 results. XXV.
Fig. 24. Line elements showing the orientation of the projected magnetic field (90◦to the synchrotron polarization angle) from the Planck 28.4 GHz
data at 2◦resolution. The region within b=±10◦has been masked out (grey rings) to avoid crowding of the vectors. The colour scale shows the
polarized intensity. The maps are in orthographic projection, centred on the Galactic poles, north (left) and south (right). Light dashed circles show
the outlines of Loops I to IV, as defined by Berkhuijsen et al. (1971) and shown in Fig. 20; the red dashed line is the locus of filament IIIs from
Vidal et al. (2015).
Fig. 25. Combined 28.4 GHz polarization map at 1◦resolution of the
North Polar spur and the region outside its perimeter, including fila-
ment X of Vidal et al. (2015), which leaves the Galactic plane at a shal-
low angle from near l=36◦, and the l=45◦feature between filament X
and the NPS. Colour scale is asinh, but the length scale for polarization
line elements remains linear, as in all our figures.
The interstellar medium within 200–300 pc has now been to-
mographically mapped in a variety of absorption tracers against
stars with Hipparcos distances (Frisch et al. 2011). Towards the
bright left side of Loop I, the LC wall is at D=70–80 pc,
and in the Galactic plane dense gas then extends continuously
to the Ophiuchus and Lupus molecular clouds at D=140–
160 pc (e.g., Vergely et al. 2010). These are the remnants of the
parent cloud(s) of the Sco-Cen association. Although there is
a small cavity surrounding the most active part of the associa-
tion (Puspitarini & Lallement 2012), there is no sign of a cav-
ity with the angular size of Loop I, and this can no longer be
ignored given that microwave polarization maps show that the
loop crosses the Galactic plane without interruption. Lallement
et al. (2014) push the analysis out to 1 kpc by using stars with
photometric distances instead of relying purely on Hipparcos.
They find a low-density cavity starting some 250 pc from the
Sun, covering the longitude range of Loop I. If this is the Loop,
the radius is roughly a kiloparsec, and it would stretch well into
the halo.
Given the lack of room for Loop I immediately outside the
Local Cavity, Frisch has long argued (e.g., Frisch 1981;Frisch
et al. 2011) that the two are continuous, and that the NPS is
merely separated from us by a fold in the LC wall. However,
this view is at odds with the evidence from soft X-ray absorp-
tion of the NPS. This absorption is visible in the ROSAT map
(Fig. 21h) where the inner boundary of the X-ray NPS at b<40◦
closely follows the border of the cold material traced in panels d–
g of Fig. 21. This material can be confidently assigned to the LC
wall, since it is picked up by starlight polarization at D>100 pc
(Santos et al. 2011), and again extinction and emission polariza-
tion is consistent. X-ray spectra in the region are well fitted by
absorbed thermal emission (Puspitarini et al. 2014). Similarly,
Sofue (2015) makes a strong case that the sharp truncation of
the X-ray NPS below b=10◦is due to absorption in the cold
clouds of the Aquila Rift. Both papers find that the hydrogen
A25, page 29 of 45
A&A 594, A25 (2016)
column density derived from the X-ray spectra requires a dis-
tance to the front of the spur, roughly the shell tangent-point,
of at least 200 pc12 . The corresponding distance to the centre is
D>400 pc.
Further evidence for a larger distance comes from the
21-cm polarization map of Wolleben et al. (2006). As noted by
Wolleben (2007), the emission from Loop I is strongly depo-
larized at |b| ≤ 30◦. Given the cut-offin latitude, which does not
correspond to any feature in the intrinsic intensity or polarization
structure, the depolarization is almost certainly caused by fluctu-
ations in the foreground Faraday depth, φF, across the 30-arcmin
beam of the Wolleben map. The sharp cutoffin bsuggests that
the high-latitude part of the loop projects above the dense layer
of the ISM that creates the “Faraday horizon”. At the near dis-
tance, the implied scale height, h≈50 pc, is much too low for the
most effective source of widespread Faraday rotation, the warm
ionized medium (WIM) (h=1–2 kpc, Gaensler et al. 2008).
The Faraday depths required also imply a larger distance.
Depolarization requires fluctuations with σF>
∼π/2λ2=
36 rad m−2. The LC has negligible Faraday depth, due to its
low density, and so the magneto-ionic medium responsible must
be in the LC wall, if Loop I is directly adjacent to us, even
though the wall is mainly traced by neutral gas. Fluctuations
across a beamwidth θrequires field tangling on scales Dθ, and
hence the Faraday depth along the line of sight of length Lcen-
tred at distance Dintegrates as a random walk, with amplitude
σRM =kFhne|Bk|i√LDθ, where kF=8.1×10−6m nT−1pc−1.
The WIM has electron density ne≈1 cm−3at pressure equi-
librium. We also assume B=1 nT, with hBkia factor 1/√3
smaller. Inserting these values and taking D=80 pc, the re-
quired path length is L>
∼70 pc, which places the front face of
the loop beyond the Sco-Cen association, even if the WIM has
unity filling factor. Of course higher-density material would re-
quire less path and in fact the Sh 2-27 H ii region (G006+236),
ionized by ζOph at D=112 pc (van Leeuwen 2007), casts a
particularly deep Faraday depolarization “shadow” at 21 cm, im-
plying it is in the foreground despite a precise distance that puts
it near the assumed shell centre. However, such objects are read-
ily detected in Hαsurveys, and cover only a small fraction of the
depolarized region. The near distance can only be made com-
patible with these results if the polarization is dominated by the
far hemisphere of the loop, as in one of the models discussed in
Sect. 5.3.3.
Given all this evidence that Loop I is several times further
away than the Sco-Cen OB association, we searched for an al-
ternative group of OB stars13 to generate the Loop. Given the
angle at which the NPS passes through the Galactic plane, these
must be at positive latitudes; we searched −45◦<l<15◦,
5◦<b<35◦. Apart from Sco-Cen, in this region there are
no known OB associations (de Zeeuw et al. 1999), but an asso-
ciation old enough to generate Loop I could already have lost its
O stars and could easily have escaped notice. Using SIMBAD
(Wenger et al. 2000) we found 51 OB stars (mostly B giants)
with parallax, π, between 1 and 3.3 mas; 14 with 0 <π<1 mas;
and 14 with π < 0. Although the individual parallaxes are barely
significant, the distribution is strongly biassed to positive values,
12 Sofue (2015) argues for a distance of >
∼500 pc, but this relies on the
assumption that the Aquila Rift is at a single distance, whereas clouds
at several different distances are involved, see the tomographic maps of
Lallement et al. (2014).
13 We follow Reed (2003) in defining OB stars as main sequence stars
of type B2 and earlier, and giants of type B9 or earlier, roughly corre-
sponding to stars massive enough to go supernova.
so there certainly has been star formation within the last 20 Myr
in the relevant volume. We will have to await Gaia data for a
definitive assessment of stellar groupings there.
5.3.3. Interpretation
The majority view of Loop I is that it is a nearby pre-existing
cavity re-energized by one or more recent supernova, in order to
explain the cold neutral border seen in dust and H i(see Salter
1983 for a review).
The H iborder appears at VLSR =0 km s−1(Fig. 21d),
which is uninformative, because expansion at the tangent points
is perpendicular to the line of sight. Heiles (1984) interpreted
the border as part of a shell expanding at ≈25 km s−1; this im-
plies that the obvious Sco-Cen supershell, visible most clearly
at −10 km s−1(e.g., Kalberla et al. 2005;Vidal et al. 2015,
Fig. 21e), is merely the end-cap of Loop I itself. It is worth
noting that this shell is elongated in the same direction as our
proposed outline for Loop I. However, the Sco-Cen shell is more
usually seen as a distinct structure with a smaller angular size
than Loop I (e.g., de Geus 1992). If Loop I has a much lower
expansion speed, e.g., 2 km s−1as estimated by Weaver (1979),
its H iemission would not be separable in velocity from the very
local H isurrounding the Sun. The difference between these two
interpretations is dynamically significant: at 25 km s−1, the cold
shell cannot have been accelerated by a shock, since the heating
would have dissociated the gas and evaporated the dust; therefore
acceleration by a pressure gradient as in the stellar wind model
proposed by Weaver (1979) is required. A 2 km s−1expansion is
consistent with the weak remains of a shock, although of course
a much faster shock would be needed to heat the interior gas
above 106K. We also note the absence of optical emission lines
characteristic of a cooling shock, which would be expected if the
shock speed exceeded 10 km s−1. A plausible explanation is that
the re-energizing blast wave has only recently hit the cavity wall,
which would also explain why the NPS is so much brighter and
more sharply-defined than the other large loops (e.g., Borken &
Iwan 1977).
Figure 24 shows the projected magnetic field orientation
around the Galactic polar caps. We show the Planck data only,
smoothed to 2◦resolution, because of the suspect large-scale
structure in the WMAP data. Figure 25 shows the pattern closer
to the Galactic plane. The quasi-parallel field pattern in the spurs
has been recently discussed by Vidal et al. (2015). Here we wish
to draw attention to the regions of organized field orientation at
high latitude, outside the NPS, which parallels it quite closely
over most of its length. Some of this region is occupied by the
cold border, where the magnetic field may be organized by the
expansion of the loop, e.g., as modelled by Vidal et al. (2015).
However, Fig. 24 shows that the parallel-field region extends past
longitudes of 60◦, well beyond the outer limit of the cold border.
The origin of this high-latitude, inter-loop polarization is un-
clear. Wolleben (2007) sees it as emission from a “new loop”,
in whose shell the solar system is embedded, but the evidence
for this loop is weak (Sect. 5.4) and its proposed interpenetra-
tion of Loop I is, to say the least, dynamically problematic. In a
somewhat similar picture, Vidal et al. (2015), following Heiles
(1998), place the Solar System on the surface of Loop I, so
that loop emission fills a hemisphere: this can explain the high-
latitude polarization, but it is not clear why the cold border seems
to have a significantly smaller radius than 90◦. If the loop is
nearby and the polarization is from the halo, the alignment with
Loop I would be an accident. However, especially if the loop is
as large as suggested in Sect. 5.3.2, the emission may be from
A25, page 30 of 45
Planck Collaboration: Planck 2015 results. XXV.
the immediate environment of the loop, suggesting that Loop I
is brightest where the interstellar field is in the plane of the shell
surface. This is exactly as predicted by Spoelstra (1972), who ap-
plied the model of van der Laan (1962) to the NPS. This appeals
to the asymmetry of magnetic stress, which has a net tension
along the field lines, so magnetic forces oppose expansion only
perpendicular to the field. The same model has been invoked to
explain the occurrence of barrel-shaped SNRs aligned with the
Galactic plane (Gaensler 1998).
Much of the random component of the Galactic magnetic
field is on scales smaller than Loop I (e.g., Haverkorn et al. 2008;
Brown & Taylor 2001), so the external field around it is likely to
show significant variation in magnitude and direction. Such large
objects are therefore unlikely to show the bilateral symmetry of
classic barrel-shaped remnants. Nevertheless the basic physics
still operates; thus, segments of the shell expanding perpendicu-
lar to the field, along the edge of the NPS, will be more strongly
impeded and so require higher internal pressure, and therefore
stronger nonthermal and thermal (X-ray) emissivity. By the same
token, we expect the fainter parts of the shell to expand fastest,
consistent with our interpretation of the bottom-right arcs as a
bulge in the loop boundary. Of course, asymmetric densities in
the ISM will also play a role, and as noted above, if the loop is
at D≈140 pc the observed density gradient is in the required
sense.
In old supernova remnants the projected magnetic field is
commonly aligned with the outer surface, i.e., the shock prop-
agating through the ISM. This alignment is naturally produced
by shock compression, which ensures that the field lines behind
the shock tend to be parallel to the shock plane. In this picture
there is no global field order in the shock plane, but the field
anisotropy causes polarization at viewing angles other than face-
on, reaching a maximum when the line-of-sight is in the shock
plane (see Laing 1980). Applied to Loop I, this model implies
that both the NPS and the internal spurs are tangential views
of shock fronts. There is no problem with this if the internal
spurs are background features, as suggested for one of the most
prominent in Sect. 5.5 below. However, if, as seems likely, at
least some of the internal spurs are physically associated with
Loop I, then we have to explain multiple concentric shock waves.
The expected supernova rate for the Sco-Cen OB association
is around 1 Myr−1(de Geus 1992), while remnant lifetimes are
conventionally .30kyr, and in the low-density environment of a
superbubble they are likely to be shorter. Therefore multiple ac-
tive supernova remants within the superbubble are very unlikely,
even if Loop I is powered by a much more massive OB associ-
ation than Sco-Cen. However, if a re-energizing blast wave has
recently run into the contact discontinuity at the outer surface of
the superbubble, it will separate into a transmitted and reflected
shock, and for any non-ideal geometry a rather complex pattern
of crossing shocks is expected in the reflected wave. MHD sim-
ulations of this scenario would be illuminating, but are beyond
the scope of this paper.
A very different model for the NPS has been proposed by
Heiles (1998) and is implicitly applied to the internal filaments
of Loop I by Vidal et al. (2015). In their geometry the NPS is
not a tangential view of the superbubble boundary, since the so-
lar system is on the bubble surface. The organized field pattern
is due to the wrapping of a relatively ordered interstellar field
over the surface of the expanding bubble. The spurs are bundles
of field lines with an enhanced density of cosmic ray electrons
and therefore enhanced synchrotron emission. An attractive fea-
ture of this is that a very simple geometric model gives a good
qualitative fit to the field pattern in the loop. However, the model
has the peculiarity that the emission is implied to be dominated
by the far side of the loop; the near side, in which the viewer
is embedded, would contribute a parallel polarization over the
whole hemisphere, quite unlike the observed pattern. We note
that in this model, Sh 2-27 can depolarize the loop emission even
though located near the centre of the bubble.
A hybrid of the two models is possible if the solar system
is moved offthe loop surface in the Heiles geometry, so that
the NPS resumes its usual role as the projected loop boundary.
However, this would worsen the agreement of the model and
observed field pattern, and is still subject to the problem that
emission from the near surface of the bubble would give a much
more uniform projected field.
5.3.4. Possible contamination of the CMB by Loop I?
Liu et al. (2014) have recently argued that emission from
Loop I is contaminating current microwave background maps.
Specifically, they demonstrate an alignment between the BHS
small circle and positive peaks in the low-multipole (`≤20)
structure of the WMAP ILC map (Bennett et al. 2013a);
von Hausegger et al. (2016) confirm that the alignment is also
seen in the Planck CMB maps14. Showed that the analysis by Liu
et al. underestimated the probability, p, of a chance alignment,
but von Hausegger et al. (2016) present a stronger statistical test
that gives p<3×10−3. The amplitude of these peaks is about
ten times larger than worst-case errors due to foreground con-
tamination in the ILC maps estimated by Bennett et al. (2013a),
which occur around the edge of the mask recommended for use
with their ILC map at low multipoles (Liu et al. applied no mask
in their analysis). The most convincing alignment with the cir-
cle is in the southern Galactic hemisphere, including along the
bottom-right section that we note above shows no sign of emis-
sion from the loop border. In general these peaks in Liu et al.’s
low-lmap show little correlation with the actual synchrotron or
dust emission from Loop I, especially if we mask the Galactic
plane as Bennett et al. recommend. We note that Liu et al. (2014)
suggest that the contamination might be due to magnetic dipole
emission from dust grains associated with Loop I. However, the
securely-detected dust associated with Loop I is mostly located
at larger radii than the synchrotron ridge, and would not itself
give a significant signal in the analysis of Liu et al. (2014) or
von Hausegger et al. (2016).
Moreover it seems unlikely that the dust in Loop I is par-
ticularly unusual; even if we posit a unique dust type in the
loops, Mertsch & Sarkar (2013) have argued convincingly that
similar structures are scattered throughout the Galactic disc and
contribute a significant fraction of its synchrotron emission. By
the same token, the proposed dust emission from distant loops
should accumulate along the Galactic plane, but no such signal
is apparent in the CMB maps. von Hausegger et al. (2016) ar-
gue that this expected signal is restricted to |b|<5◦, a region
where component separation is unreliable (hence the need for a
Galactic mask). While it is true that it is not possible to accu-
rately estimate the CMB signal in this region, we can rule out
foreground contamination at the level required by this model.
The integrated brightness of the distant loops is is an order of
magnitude larger than that of Loop I, and the features respon-
sible for the Loop I alignment are among the brightest in the
CMB maps after filtering to l≤20, with amplitudes 50–100 µK.
Using this to scale the arbitrary units in Fig. 5 of von Hausegger
et al., the brightness from distant loops on the Galactic plane
14 Ogburn (2014).
A25, page 31 of 45
A&A 594, A25 (2016)
Fig. 26. Mean brightness vs. Galactic latitude, b, in the Planck CMB
maps from Planck Collaboration IX (2016), in the range |l|<60◦. Solid
green: NILC; dotted blue: SMICA; dot-dashed magenta: SEVEM (all at
5 arcmin resolution); dashed red: baseline Commander model at 1◦res-
olution (Planck Collaboration X 2016). The first three methods all use
inpainting to fill masked regions around bright compact features in some
or all of the input maps: pixels affected by this have been omitted from
the averaging. High amplitudes close to b=0 are due to residual fore-
ground contamination.
averaged over |l|<60◦should be ≈900 µK, and is still over
100 µK at |b|=4◦. Figure 26 shows profiles of derived CMB
brightness against bfor the four main Planck component separa-
tion methods. Except for SEVEM, which is the method least well
adapted to removing diffuse foregrounds, substantial contamina-
tion from Galactic foregrounds is confined to |b|<2◦. Even the
SEVEM map has smaller Galactic residuals than predicted from
the von Hausegger et al. analysis.
In short, there is little reason to believe in the contamination
suggested by Liu et al. The alignment they suggest is with a theo-
retical schematic rather than the observed synchrotron loop, and
is even further away from the actual dust emission in the loop.
The formal improbability that the observed alignment occurs by
chance, p≈0.2% in the best (i.e., Planck) CMB maps, should
be offset against the low prior probability that this schematic ac-
cidentally aligns much better with a previously-undetected struc-
ture than with the synchrotron loop that it was intended to model.
This p-value does not allow for any look-elsewhere effect such
as potential alignments with other loops. Moreover, their pro-
posed physical model implies additional strong contamination
along the Galactic plane, which is not observed.
5.4. Other loops and spurs
Three other synchrotron loops that are generally agreed upon
(e.g., Vidal et al. 2015), and there are a dozen or so shorter fil-
aments or spurs, with several mutually-incompatible proposals
for joining some of them into loops. Given our emphasis on the
non-circularity of Loop I, and the fact that it seems all too easy to
find small circles passing plausibly close to a sequence of ridges
on the sky maps, we do not find any of these newly proposed
loops convincing; however, their component spurs are certainly
real. We discuss the more interesting examples in turn.
Loop II. Also known as the Cetus arc, Loop II is a very dif-
fuse, 45◦-radius structure centred at (l,b)=(100◦,−32◦
.5). The
faint, broad emission on the left-hand side of the maps in Fig. 2
is attributed to it. As originally defined by Large et al. (1962),
the right-hand end of the arc was considered to be the “South
Polar Spur” (SPS), which descends from the plane at l=45◦
(visible in Figs. 20 and 21). The SPS is prominent in both the
unsharp masked 408 MHz map and the synchrotron polarization
map, but curves towards the right as it leaves the plane (as does
its projected magnetic field), in the opposite direction from the
notional path of Loop II; it therefore appears to be unrelated. The
BHS small circle for Loop II just crosses the plane, but no emis-
sion from it has been traced in the northern Galactic hemisphere;
the top of the loop may be obscured by Galactic plane emission,
but it is also possible that the expanding shell was halted by the
dense gas in the plane. Alternatively, Weaver (1979) and Heiles
(1998) suggest that Loop III is its northern extension, although
the two loops are offset by 24◦in longitude. No reliable distance
information is available for Loop II.
Figure 24 shows a coherent field pattern roughly parallel to
the locus of Loop II that is not visible in the WMAP data. The
polarized intensity essentially disappears into the noise at the
bottom of the loop near the South Galactic Pole (SGP), but a
coherent polarization field can be traced along a path that goes
a few degrees further towards the SGP than the BHS fit, and
closes a few degrees higher in longitude from the SPS. Some of
this polarization is detected at 1.4 GHz by Wolleben et al. (2006),
but the polarization angles are more disordered, presumably due
to Faraday rotation, and so the alignment with the loop is not so
clear.
Loop III. Loop III is centred at (l,b)≈(124◦,15.
◦5); BHS give
a radius of 32.
◦5, although Figs. 20 and 24 show that polar-
ized emission is detected several degrees further out. Loop III
is clearly detected around most of its perimeter in the north-
ern hemisphere in both total intensity and polarization. Its right-
hand spur, rising from the plane directly above the Cyg X re-
gion, passes through the region of deep coverage near the North
Ecliptic Pole (l,b)≈(96◦,30◦) on the Planck maps, giving a par-
ticularly clean view of the well ordered magnetic field (Fig. 27).
This Cygnus spur is relatively narrow (FW HM ≈5◦) until it
reaches b≈30◦after which the loop becomes much broader and
fainter. As it returns to the plane its polarization is obscured by
the strongly polarized Fan region (100◦<l<170◦), in which
the magnetic field is parallel to the plane and therefore orthog-
onal to the expected field in the loop. BHS argue that the loop
re-emerges from the fan region in the southern hemisphere at
l≈150◦, and the B-field pattern in our synchrotron polarization
map is consistent with this, but the loop does not convincingly
close in the south. Vidal et al. (2015) imply that their filament
IIIs may be associated with Loop III proper, but this is one of
the regions where WMAP and Planck disagree the most, and the
existence of this filament is not confirmed in the Planck maps,
except for a short section near (l,b)=(82◦,−30◦) (see Fig. 24).
Kun (2007) discusses possible interactions between Loop III and
the cold gas north of the plane, notably the North Celestial Loop
(Meyerdierks et al. 1991) that may be swept-up material.
Loop IV. Loop IV is centred at (l,b)=(315◦,48.5◦) and is
projected entirely inside Loop I. Only an arc subtending ≈20◦
along its low-latitude rim is clear in our polarization map; see
Vidal et al. (2015) for analysis. The high-latitude rim of Loop IV
that parallels the top of the NPS is very noisy here but clear in
the 21 cm polarization map of Wolleben et al. (2006).
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Planck Collaboration: Planck 2015 results. XXV.
Fig. 27. Combined polarization map of the Cygnus spur, the right-hand
end of Loop III. The map is centred at (l,b)=(90◦,+27◦) and the gratic-
ule has a 15◦spacing.
The South Polar Spur (SPS) and Filament X. As discussed
above, the SPS was once seen as the low-longitude edge of
Loop II, but in fact curves in the opposite direction so the outside
edge of the SPS is to its right. It also forms the most prominent
segment of the S1 loop proposed by Wolleben (2007), and was
labelled filament VIIb by Vidal et al. (2015). Wolleben’s view
of the SPS was obscured by Faraday rotation near the plane; as
noted by Vidal et al. (2015), the high frequency (&20 GHz) po-
larization structure of the SPS is more strongly curved near the
plane than expected for Wolleben’s path. The other section iden-
tified by Wolleben as part of his new loop is usually seen as one
of the interior ridges of Loop I and the northern part of Loop IV;
thus, all components of loop S1 appear to be parts of smaller
structures and there is no clear evidence that S1 exists as a co-
herent physical structure.
The SPS has its own cold border just outside the radio ridge,
at l≈60◦,−30◦>b>−70◦, which is particularly clear in dust
polarization (Fig. 28). The border is visible in H iat 0 < vLSR <
10 km s−1(Fig. 21d,e), which implies a distance of less than a
few hundred parsecs.
Given our discussion above, we are reluctant to define a
loop by extending a small circle from the clearly-detected arc
of the SPS. In fact, a slightly fainter spur immediately north of
the plane (Vidal et al.’s filament X) seems to be a reflection-
symmetric twin of the SPS, suggesting a wasp-waisted cavity
shaped by the dense gas in the plane (see the left-hand edge of
Fig. 21c).
Our map allows filament X to be traced further north, to
b=37◦, where it curves over and enters the region of bright
diffuse polarized emission just outside the NPS (Fig. 25).
Unusually, the field in this northern section is perpendicular to
the filament. This alignment is not due to a contribution from
the underlying diffuse emission; if it were, the filament would
appear as a trough (due to cancellation), not a ridge in polarized
intensity. A possible return section is marked in Fig. 20.
Fig. 28. The South Polar Spur. Left:Planck-WMAP polarization map;
right:Planck 353-GHz dust polarization. The maps are centred at
(l,b)=(50◦,−45◦) and the graticule has a 10◦spacing.
The Orion-Taurus Ridge and Filament XI. Unlike the other
spurs, these features do not emerge from the Galactic plane but
run parallel to it at closest approach, at (l,b)≈(220◦,−15◦), and
(l,b)≈(230◦,40◦), respectively (see Fig. 20). They are visible
in both unsharp-masked total intensity and in polarization (see
Fig. 2 of Vidal et al. 2015). Milogradov-Turin & Uroševi´
c(1997)
and Borka et al. (2008) propose the existence of a pair of very
large, overlapping loops (V and VI in their designation) that con-
nect a sparse collection of features including the right-hand end
(at l<180◦) of the Ori-Tau ridge and Vidal et al.’s filament IIIs.
However, the rest of the Ori-Tau ridge curves far away from the
locus of Loops V/VI, and the magnetic field continuously fol-
lows that curvature, so these loops appear to be spurious.
Filament XI is an approximately 80◦-radius arc that seems to
curve away from the plane at both ends (Fig. 20); however, it is
unclear if features near l=264◦belong to the left-hand end of
Filament XI or the right-hand side of Loop I (or filament XII, if
not part of Loop I). Extended to a full loop, Filament XI encloses
most of Galactic quadrant 3 north of the Plane, which is one of
the faintest regions of the polarized sky; over much of this quad-
rant no polarization is clearly detected by Planck and/or WMAP,
even at 4◦resolution.
l=45◦feature. This is the highly-polarized patch covering
12◦<b<45◦and from the edge of the NPS to l≈50◦(Fig. 25).
The low-bedge of the polarization feature coincides with a rapid
brightening of the inner halo in total intensity, and when polar-
ization re-appears at b≈8◦it has an orthogonal orientation,
roughly parallel to the Galactic plane. Hence this minimum is
due to a cancellation between the plane-parallel field in the in-
ner halo and the near-vertical field in our feature, and does not
mark the physical edge. A faint ridge runs through the top half
of the feature at l≈45◦, which is just visible in the 408 MHz
map (Fig. 21a,b); otherwise there is little trace in total intensity.
The magnetic field runs at a slight angle to the ridge of the NPS
throughout; in the adjacent section of the NPS the projected field
A25, page 33 of 45
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is misaligned with the ridge-line and in better agreement with
that in the l=45◦feature. Plausibly the feature overlaps in pro-
jection with Loop I, but if so the current data do not allow us to
trace it inside the loop.
Smaller loops. Mertsch & Sarkar (2013) have emphasized that
analogues to the large synchrotron loops should be common, and
more distant and smaller examples should therefore be visible in
the sky maps close to the Galactic plane. One rather clear ex-
ample is the polarization structure south of the Cyg X region,
centred at (l,b)=(90◦,−5◦), which seems to be a superposition
of a 5◦-radius loop around this bright star-forming complex and
another spur heading to the lower right (Fig. 29). The radius is
about 120 pc, assuming a distance of 1.4 kpc for Cyg X (Rygl
et al. 2012). A northern counterpart is less obvious, partly be-
cause of the two bright supernova remnants, HB 21 and W 63,
that lie on the likely path. Figure 29 also shows another possible
distant loop: the set of three arches in the southern hemisphere,
centred near l=60◦; they are also marked in Fig. 20. The coher-
ence of the polarization vectors allows the outer arch to be fol-
lowed around to (l,b)=(17◦
.4,−12◦
.9), just south of the Cygnus
Loop supernova remnant. While the polarized intensity suggests
that the inner arch returns to the plane at l≈61◦, the polariza-
tion vectors are orthogonal to the apparent ridge-line, and so this
might be a different structure. Figure 29 also shows that the dust
polarization has the same magnetic pattern as the synchrotron
emission around Cyg X. There is also a polarized dust filament
that runs roughly along the inner edge of the middle l=60◦arch,
while the outer arch parallels the dust in polarization as it runs
behind the SPS.
5.5. Fermi bubbles in polarization
The Fermi bubbles are two large structures extending perpen-
dicular to the Galactic plane, up to about ±55◦. They were
first discovered by Dobler et al. (2010) using spatial templates
while searching for a γ-ray counterpart to the microwave haze
(e.g., Finkbeiner 2004;Planck Collaboration Int. IX 2013) in the
Fermi data.
The origin of the bubbles is not clear, although their location
and symmetry with respect to the Galactic centre suggest that
they originate there. They have a γ-ray spectrum significantly
harder than the inverse Comptom emission from the Galactic
Halo or the one from pion decays from collisions of CR, with
the ISM protons and heavier nuclei (Su et al. 2010). Different
models have been proposed to explain their origin, most of them
relating to a recent AGN-type activity at the Galactic centre. For
a recent overview on the bubbles and their spectral behaviour in
γ-ray, see Dobler (2012) and Ackermann et al. (2014).
We compare our combined Planck-WMAP polarization am-
plitude map at 28.4 GHz with a full-sky Fermi map from
Ackermann et al. (2014). The γ-ray map was produced using
50 months of Fermi LAT data (Atwood et al. 2009), using only
the “UltraClean” class events, i.e., the sample of events with less
contamination from misclassified interactions in the Fermi LAT
instrument. See Ackermann et al. (2014) for a detailed descrip-
tion on the Fermi data analysis. We use the high energy map,
which covers 10–500 GeV. In this energy range, the emission
from the Fermi bubbles appears clearly without any component
separation at latitudes b&20◦. The main foreground emission
in gamma rays is from decay of π0particles produced by CR
protons interacting with the ISM (Ackermann et al. 2012). If the
CR density and spectrum are roughly spatially uniform over the
Fig. 29. Top: combined polarization map with magnetic field vec-
tors overlaid, centred on (l,b)=(70◦,−7◦
.5) showing several polar-
ized arches south of the plane. Bright compact features along the
plane are all SNRs: DA 530 (G93.3+6.9); HB 21 (G89.2+4.7); W 63
(G82.2+5.3); CTB 87 (G74.9+1.2); the Cygnus Loop (G74.0−8.5);
CTB 80 (G69.0+2.7); and W 51C (G49.1−0.6). The large spur at l≈
50◦is the SPS. Bottom: dust polarization in the same region as seen
at 353 GHz. Maps are rectangular projections and have an asinh colour
scale.
Galaxy, then the π0emission will be proportional to the ISM
column density (Su et al. 2010). A good tracer for the ISM col-
umn density are maps of thermal emission from dust grains,
since dust is well mixed in the ISM and its emission is opti-
cally thin. Here, we use the Planck 353 GHz optical depth map
from Planck Collaboration XI (2014) as a column-density proxy
for π0emission to fit and remove it from the γ-ray map.
In the middle panel of Fig. 30, we show the resulting 10–
500 GeV Fermi map. A region of about ±5◦along the Galactic
plane, where the subtraction of the scaled Planck 353 GHz map
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Planck Collaboration: Planck 2015 results. XXV.
Fig. 30. Left: cartographic projection of the combined Planck-WMAP polarization amplitude, centred at the Galactic centre, with graticule lines at
15◦intervals. The regions defined by red and yellow dashed lines were selected to calculate a polarized spectral index; in red is the filament around
the Fermi bubble and in yellow a control area. Centre: cartographic projection of the 10–500GeV map from Ackermann et al. (2014), with the
π0emission subtracted as described in the text, showing the Fermi bubbles. Right:Fermi bubbles component from Selig et al. (2015). The black
outline corresponds to the centre of the narrow filaments visible in the polarization map on the left.
left large residuals, has been masked out and is shown in grey.
The left panel in Fig. 30 shows our combined Planck-WMAP
polarization amplitude map, where narrow filaments are visi-
ble that correspond to the border of the Fermi bubbles. We note
that the northern filament is much clearer than the southern one.
The ridge-lines of these filaments is over-plotted in the Fermi
map on the right to show the correspondence. It is remarkable to
see how closely the polarized filament follows the border of the
Fermi bubbles. We also show on the right panel the denoised
and source-subtracted Fermi bubbles template constructed by
Selig et al. (2015) using a Bayesian inference algorithm. Their
reconstruction has a poorer angular resolution than the less pro-
cessed version we show at the centre of Fig. 30, but still it is
clear that the synchrotron filaments outline the border of the
bubbles. The polarized filaments are unresolved in the Planck
polarization maps (400FWHM beam), implying that these are
very narrow synchrotron structures. This is further confirmation
of a sharp border for the Fermi bubbles, and models that fail to
predict this feature are now less favoured (see Dobler 2012 for a
description of some of these models).
Polarization spectral index. We have measured the polarization
spectral index, βpol, between the bias-corrected WMAP 23 GHz
and Planck 30 GHz using T–Tplots of a region that encom-
passes the north filament. The left panel of Fig. 30 shows two
regions in the polarization intensity map that we selected to
measure βpol, one including most of the northern filament (in
red) and a second region that we use as control (yellow).
The T–Tplots are shown in Fig. 31. The polarization spec-
tral index of the Fermi bubble filament is βpol =−2.54 ±0.16,
while the nearby control region has βpol =−2.90 ±0.36. The
spectral index of the region that includes the filament is much
flatter than the standard value of −3.0 normally assumed for dif-
fuse synchrotron emission. It is also flatter than all the values
presented in Fuskeland et al. (2014) and Vidal et al. (2015), who
explored the polarized spectral indices in different regions us-
ing WMAP data. This flatter spectral index indicates that the
energy distribution of the synchrotron radiation along the fil-
ament has more higher energy electrons than that of the dif-
fuse component next to it. The polarization spectral index of
the Fermi bubble filament we measure here is identical to the
spectral index of the microwave haze of βHaze =2.54 ±0.05
(Planck Collaboration Int. IX 2013). This supports the relation-
ship between the filament and the haze/bubble emission.
Fermi bubbles and Loop I. A relationship between Loop I (see
Sect. 5.3.1) and the Fermi bubbles has been hypothesized (e.g.,
Kataoka et al. 2013). This would put Loop I at the Galactic cen-
tre, with a much larger size and radio luminosity, as long argued
by Sofue (e.g., Sofue 1977,1994,2015).
With the new Planck-WMAP combined polarization map,
in Sect. 5.3.1 we trace Loop I below the Galactic plane. It is
A25, page 35 of 45
A&A 594, A25 (2016)
Fig. 31. T–Tplots between the polarization amplitude maps of WMAP
K-band and Planck 28.4 GHz. The two plots correspond to the regions
defined in the left panel of Fig. 30. The Fermi bubble filament shows a
flatter spectrum than the control region.
clear from the top panel of Fig. 20 and from Fig. 30 that the
southern part of the Fermi bubble extends outside Loop I, and
there is no trace of an interaction with the bubble in the radio
maps, given the continuity and smoothness of the Loop I polar-
ized filaments in the southern hemisphere. The two structures
must therefore be at different distances. The Fermi bubbles are
well centred on the Galactic centre, show a pinched structure
symmetric about the Galactic plane, as expected for an outburst
from the nucleus (Sofue 1994), and are unique in the γ-ray sky.
All this makes it highly likely that they are located at the dis-
tance of the Galactic centre, as usually assumed. In contrast,
Loop I is centred 35◦from the nucleus and significantly above
the plane, while on its left-hand side, where it can be clearly fol-
lowed, it appears to extend through the plane without any sign
of deviation (Sect. 5.3.1). This is just as expected if it is rela-
tively nearby and embedded in the disc rather than extending
far beyond it. Of course, it also has a larger angular size than
the Fermi bubbles, and there are at least several similar loops
(Sect. 5.4). We are therefore confident that Loop I is a foreground
feature.
Fig. 32. Top left: Hαmap in the velocity range −80 <Vlsr <−40 km s−1.
Notice the vertical filament that runs at l≈75◦and −60◦<b<−20◦.
Top right: Faraday depth map from Oppermann et al. (2012), which also
shows a filament at the same location, with a mean value of −25 rad m−2
along its extension. The filament has a counterpart in “absorption,” vis-
ible as a trough in polarization intensity maps. Bottom left: WMAP
23 GHz polarization intensity map. Bottom right:Planck-WMAP polar-
ization intensity map. The Hαmap on the top left corner has an angular
resolution of 1◦, while the other three maps have a common resolution
of 3◦. The grid spacing is 15◦
.
5.6. An anti-correlation of an Hαfilament and polarized
intensity
Another interesting aspect that can help in the determination of
the distance to some of the polarized structures is the discovery
of a region that shows an anti-correlation between Hαemission
and the polarized intensity in our combined Planck-WMAP map.
This has not been noticed before as far as we are aware, because
the filaments are located at high Galactic latitude, in the Faraday-
thin regime.
In Fig. 32 we show an Hαmap, integrated over the velocity
range −80 <Vlsr <−40 km s−1, from the Wisconsin H-Alpha
Mapper (WHAM) survey (Haffner et al. 2003). The filament that
runs along l≈75◦is about 40◦in length. The WMAP 22.8 GHz
polarized intensity, and the Planck-WMAP polarized intensity
map also shown in Fig. 32, show a trough at the same location as
the Hαfilament. This feature is also visible in the Faraday depth
map from Oppermann et al. (2012), shown on the same figure.
The Faraday depth, φis proportional to the magnetic field along
the line of sight and the electron density:
φ=kFZ0
r0
dr ne(r)Br(r).(5)
We have ruled out a chance correlation since the polarized fea-
ture is also clearly visible, at higher angular resolution, in the
maps from the GALFACTS (Taylor & Salter 2010) consortium
A25, page 36 of 45
Planck Collaboration: Planck 2015 results. XXV.
(Jereon Stil, priv. comm.), and the correlation is still strong at
angular scales of a few arcmin.
The origin of the observed anti-correlation between Hαin-
tensity and polarization amplitude is not clear. A first possibility
is that the trough visible in the combined Planck-WMAP po-
larization map is a depolarized region, meaning that the Hαfil-
ament lies in between the polarized background emission and
us. If this is the case, the ionized gas traced by the Hαmap
could produce Faraday rotation, due to the presence of a mag-
netic field in the plasma, depolarizing the diffuse background
emission along its extension. However, this hypothesis is not
compatible with the low density of the ionized gas. The inten-
sity of the filament in Hαis 2 R above the background at the
original resolution (60) of the Finkbeiner (2003) Hαmap. This
corresponds to a mean electron density of 2.0 cm−3(assuming
that the filament has a thickness of 1 pc), which is very low to
produce significant Faraday rotation at 23 GHz for typical val-
ues of the magnetic field. Moreover, the Faraday depth map from
Oppermann et al. (2012) has a value of about 25 rad m−2along
the filament, which corresponds to a 0◦
.3 change in polarization
angle at 23 GHz. Therefore, Faraday rotation of this high lati-
tude filament is not enough to cause any major depolarization at
23 GHz.
A second alternative is that there could be a strong coherent
magnetic field parallel to the line of sight along the filament.
This might be the result of the dynamical processes that created
the filament. If this is the case, there would be less polarized
emission at the filament location than around it.
Lastly, it might well be that the synchrotron emission from
the region of the Hαfilament is intrinsically weakly polarized.
This might be due to a less organized magnetic field in this re-
gion in comparison with the diffuse emission seen in the vicinity
of the filament on the polarization map.
We also note the fact that the Hαfilament is only visible at
negative radial velocities. This corresponds to the velocity range
of the Perseus arm of the Galaxy. If the Hαfilament belongs
to that arm, it would imply that the distance to the diffuse syn-
chrotron background is much larger than a few hundred parsecs,
lying at least 2 kpc away from us.
5.7. Limits on AME polarization
If AME is solely due to spinning dust particles, then we ex-
pect it to have a very low polarization percentage. The level
of spinning dust polarization depends on the alignment effi-
ciency of the small grains and PAHs in the interstellar magnetic
field. Lazarian & Draine (2000) considered resonance paramag-
netic relaxation, which predicts .1.5% polarization for frequen-
cies &20 GHz; Hoang et al. (2013) use constraints on the align-
ment of grains seen in ultraviolet polarization to predict that
AME will have a polarization of .0.9% at frequencies above
20 GHz. At lower frequencies (.10 GHz), the polarization frac-
tion will be higher, but as the AME spectrum steeply decreases
at low frequencies while other foreground components are in-
creasing, this will be more difficult to detect than at the peak of
the AME spectrum. Given this, and the fact that the observed
polarization percentage would naturally be less than this due to
beam and line-of-sight depolarization, we do not expect to detect
significant AME polarization with Planck, i.e., the polarization
percentage should be .1%.
Observational constraints on AME polarization have so far
been placed using relatively compact, isolated clouds, where
AME is known to be strong. Rubiño-Martín et al. (2012a) re-
view the available constraints and their relation to theoretical
models. The best constraint comes from the Perseus region,
which has significant AME emission and relatively little contam-
inating synchrotron emission. Battistelli et al. (2006) reported
a weak detection in Perseus at 3.4+1.5
−1.9% at 11 GHz, while later
measurements obtained by López-Caraballo et al. (2011) using
WMAP data found a 2σlimit of <1%; Dickinson et al. (2011)
have also measured a 2σlimit of <1.4% for Perseus (as well
as an upper limit of <1.7% in ρOphiuchus), and Génova-
Santos et al. (2015) found a 2σlimit of <2.8% in Perseus at
19 GHz. Recently, Battistelli et al. (2015) have claimed to de-
tect polarized AME emission at 21.5 GHz from RCW 175, an
Hii region, where they measured a polarization percentage of
(2.2±0.2 (random) ±0.3 (systematic))% at 21.5 GHz. However,
Battistelli et al. (2015) argue that a large fraction of this sig-
nal could be residual synchrotron radiation, leaving little or no
polarized AME. There have been very few attempts to con-
strain the polarization from the large-scale diffuse AME; Kogut
et al. (2007) find that AME accounts for less than 1% of the
total polarized signal, while Macellari et al. (2011) placed a
limit of <5% based on template fitting of WMAP data. Planck
Collaboration Int. XXII (2015) placed upper limits by corre-
lating WMAP and Planck polarization maps with the Planck
353 GHz dust polarization map. They found that the correlation
turned upwards at low frequencies, and at WMAP K-band they
placed an upper limit on AME polarization of 16%; they also
noted that the upturn could be explained by dust-correlated syn-
chrotron polarization rather than AME polarization.
We use the debiased polarization map assembled using a
weighted average of both Planck and WMAP data (see Sect. 5.1)
to look for polarized emission that is correlated with diffuse
AME from the Commander solution. We note that the weighted
polarized map assumes a spectral index of β=−3.0, rather than
an AME-like spectrum. We create T–Tplots between the AME
and polarization intensities, both rescaled to µK at 22.8 GHz, in
order to measure the percentage polarization. Assuming that any
potential AME polarization has a constant percentage across the
region, then this should result in a linear correlation of the data
points in the T–Tplots. We use the uncertainty maps and assume
a model uncertainty of 10% for both the Commander AME solu-
tion and the weighted polarization map; the latter uncertainty is
subdominant in these regions. We fit for the slope of the cor-
relation, which directly gives the percentage polarization. We
repixelize the maps to Nside =64 to ensure that the pixels are
independent. We note that the debiasing method used will not
perfectly remove all of the noise bias, particularly in regions of
low S/N, but any residual bias should be lower than the thermal
noise limits where the S/N is low.
We focus on two high latitude regions from Sect. 4.2, namely
Perseus and the Pegasus plume. We show the AME and debi-
ased polarization maps in Fig. 33, and the AME–polarization
T–Tplot for each region in Fig. 34. AME dominates the total
emission at around 20–30 GHz in both of these regions. Unlike
regions such as Corona Australis and Musca, which are strongly
contaminated by highly polarized Galactic synchrotron struc-
tures (see Sect. 5.2), Perseus and the Pegasus plume are in parts
of the sky with relatively little synchrotron polarization, although
the Pegasus plume does have two polarized synchrotron features
crossing the region; these can be seen in Fig. 33 and are marked
on Fig. 20. One of these synchrotron arcs is somewhat aligned
with the AME plume, however it is unrelated to the plume,
since it extends beyond the plume towards the Galactic plane.
The morphology of the polarized and total intensity synchrotron
emission is quite different in the region of the plume, and the
aligned arc is either highly polarized or has a flat spectrum,
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A&A 594, A25 (2016)
Fig. 33. Top:Commander synchrotron component maps; middle:
Commander AME component maps; and bottom: debiased weighted po-
larization maps. Left: Perseus; right: the Pegasus plume. The colour
scales are linear here. The regions are defined and masked as per
Sect. 4.2. The Perseus region does not contain significant polarized syn-
chrotron emission, while the Pegasus plume contains several polarized
synchrotron arcs.
since it is faint in the total intensity map. Unlike the Perseus and
Pegasus plume regions, the Orion and λOrionis regions contain
strong free-free emission: while free-free emission is unpolar-
ized there may be residual leakage from total intensity to po-
larization. In addition there is a broad background of polarized
synchrotron emission in these regions, and due to their proxim-
ity to the Galactic plane there will be a correlation between syn-
chrotron and thermal dust emission, simply because both com-
ponents will be brighter towards the Galactic plane than they are
off-plane. We apply the point source mask described in Sect. 4.2:
in the regions of interest this only consists of masking one point
source near the Pegasus plume.
We find that the strongest constraint comes from Perseus,
where we measure an AME polarization percentage (compared
with the overall polarization intensity) of 0.6±0.5%, which gives
a 2σlimit of <1.6%, matching previous constraints. The Pearson
correlation coefficient is 0.20. An improved Commander estimate
Fig. 34. T–Tplots between the polarization and AME intensities (both
in units of µK at 22.8 GHz). Top: Perseus; bottom: Pegasus plume. The
scatter is mostly due to thermal noise in the polarization map (uncertain-
ties in the polarization intensity in these regions are typically around
5µK). The fitted slope gives the polarization fraction for the region,
based on the linear fit (dashed line).
of the AME in this region (see Sect. 4.2) would decrease this up-
per limit by around 40%. In the Pegasus plume region we find
a polarization percentage of 9.0±1.9%, with a Pearson corre-
lation coefficient of 0.29, which, given the presence of corre-
lated polarized synchrotron emission, we interpret to be a 2σ
limit of <12.8%. We attempt to subtract some of the polarized
synchrotron emission by assuming that the intensity is 10% po-
larized, which removes some of the synchrotron emission from
the region near the point source. This reduces the polarization
percentage relative to the AME to 5.2±1.6%, or a 2σlimit of
<8.4%, with a Pearson correlation coefficient of 0.24; assum-
ing a higher level of synchrotron polarization results in signif-
icant negative pixels in the polarization map. We also compare
the weighted low-frequency polarization map with the polarized
dust emission at 353 GHz by separately comparing the Qand
Umaps at the two frequencies, at Nside =64, in the regions of
Perseus and the Pegasus plume. Similar analysis was performed
by Planck Collaboration Int. XXII (2015) across the whole sky.
We find similar results, with no significant detection in Perseus,
A25, page 38 of 45
Planck Collaboration: Planck 2015 results. XXV.
and a weak detection in the Pegasus plume, which we ascribe
to dust-correlated polarized synchrotron emission, as described
above.
In summary, we do not find strong evidence for diffuse polar-
ized AME in WMAP/Planck data, with a 2σupper limit of 1.6%
in the Perseus region. In order to improve on current constraints,
we will need either reduced noise levels in regions that are free of
polarized synchrotron emission (such as Perseus) or the ability to
subtract accurately the synchrotron emission from less clear re-
gions. The forthcoming map of polarized synchrotron emission
at 5 GHz by the C-Band All Sky Survey (King et al. 2014) will
assist with the latter issue.
6. Conclusions
In this paper we have discussed the Galactic foreground emis-
sion observed by WMAP/Planck between 20 and 100 GHz.
The total intensity emission comprises at least three distinct
components: synchrotron; free-free; and AME, commonly at-
tributed to spinning dust. These components emit smooth con-
tinua that are typically decreasing with frequency. It is therefore
extremely difficult to separate these components with any pre-
cision. Furthermore, these individual components are often spa-
tially correlated (e.g., towards the Galactic plane).
We have applied an internal linear combination technique to
WMAP/Planck data to produce a map that is free of CMB, free-
free, and thermal dust. This has given us some insight into the
spectral form of emission at 23 GHz. For most of the intensity
analysis we rely on the component-separation products from the
Commander code. This involved fitting a parametric model to
the data. However, a number of simplifications have had to be
made, given the limited number of frequencies and complex-
ity of the foreground emission. In particular, we have not so
far been able to produce a detailed synchrotron map in inten-
sity. Instead we have had to rely on a GALPROP spectral model
to take into account the bulk of the synchrotron emission above
20 GHz. In polarization the situation is somewhat clearer. The
polarization maps at 20–50 GHz are dominated by synchrotron
radiation; free-free and AME are expected to have very low po-
larization fractions.
Nevertheless, in this study we have been able to produce rea-
sonable products for free-free and AME, from which we have
inferred some physical results. Our polarization results do not
rely directly on component separation.
The main conclusions from our study are as follows.
–After applying an ILC procedure to remove CMB, free-
free, and thermal dust emission at 22.8 GHz, we find that
the residual emission has significant spectral variations. This
emission is expected to be dominated by (weak) synchrotron
and (stronger) AME emission. In particular, we find that for
some regions the spectrum is flatter than that of free-free
emission (β=−2.1).
–The aforementioned areas are largely towards bright H ii re-
gions such as the ζOphiuchi cloud (Sh 2-27)and the
California nebula. Because synchrotron emission is not ob-
served to be flatter than about β=−2.3, the spectral changes
are thought to be due to a high frequency component of
spinning dust. Indeed, Planck Collaboration Int. XV (2014)
found evidence for a spinning dust peak at around 50 GHz
for the California nebula.
–Our free-free map appears to correspond well to previous
maps made with WMAP data, and also to Hαmaps at high
latitudes (where dust absorption is small). The amplitude
relativ to Hαis lower than the theoretical value for typical
electron temperatures, as has been found in previous studies.
–A natural explanation for the low free-free-to-Hαratio is that
there is significant scattered Hαlight. We have attempted to
estimate this fraction by correlating Hαresiduals (after sub-
tracting the mean correlation) with the dust optical depth.
We do indeed find a correlation with the thermal dust optical
depth. If we assume that electron temperature varies by no
more than ±1000 K, we find that the scattered Hαemission
accounts for 28 ±12%, or 36 ±12% after making a nominal
correction for absorption by dust.
–We also compared the Commander free-free map to the RRL
Galactic plane survey of Alves et al. (2015). Again, we find
a good overall correspondence between the maps. However,
there are significant differences, which might be due to vari-
ations in electron temperature.
–The Commander AME map is closely correlated with ther-
mal dust emission, as traced by FIR/sub-mm maps. We have
identified several new diffuse regions away from the Galactic
plane that appear to emit significant AME. These regions
are typically associated with large dust cloud complexes.
We have compared the emissivity of AME in these regions
relative to the amount of thermal dust emission. The AME
emissivity against 545 GHz and τ353 varies by a factor of
approximately 2 from region to region. The λOrionis re-
gion has a particularly high emissivity that warrants fur-
ther investigation, ideally in conjunction with additional low-
frequency data to reduce leakage between the free-free and
AME components.
–Comparison of the Commander foreground products with
similar products from the WMAP team shows that the am-
plitude of synchrotron emission is significantly lower in our
analysis. This is particularly the case at low latitudes, where
our fixed spectral model cannot account for the flatter in-
dex reported by WMAP and in earlier Planck analyses. To
compensate, our AME amplitude is significantly higher by
a factor of around 2–4 (depending on which model is be-
ing compared). The free-free component is in better agree-
ment in general, but our model is typically lower by about
10–30%, which has been noted before (Alves et al. 2010,
2012). We emphasize that our model fits the data extremely
accurately, with median residuals along the Galactic plane
of <1% for Planck and <2% for WMAP channels (Planck
Collaboration X 2016).
–We have combined WMAP/Planck polarization maps to con-
struct a higher S/N polarized intensity map. This is the high-
est S/N synchrotron polarization map available above a few
GHz, where Faraday rotation can be safely ignored.
–The new polarization map shows a number of large-scale po-
larization structures, including spurs and filaments. These
new features are detected clearly only in polarization, in-
dicating a high polarization fraction similar to the brighter
well-known spurs.
–The halo of synchrotron emission in the inner Galaxy clearly
visible at 408 MHz appears to have a low polarization
fraction (<10%) compared to the high-latitude spurs with
typical polarization fractions of 30–50%. The diffuse high-
latitude emission is also weakly (though detectably) polar-
ized (<15%).
–We discuss in detail the characteristics of Loop I. The new
polarization maps enables us to follow it further into the
southern Galactic hemisphere, where there is a substantial
deviation from the small circle fitted to the brighter northern
part, presumably due to the inhomogenous environment.
A25, page 39 of 45
A&A 594, A25 (2016)
Radio depolarization, X-ray absorption, and tomographic
mapping of the local ISM all suggest that the loop is at least
several times more distant than the traditional 120 pc.
–The projected magnetic field outside the North Polar spur
parallels the loop, supporting the model that the loop is in
part shaped by its interaction with the ambient magnetic
field.
–The filaments projected inside Loop I are mostly associated
with it, possibly a system of internal shock waves created
when a re-energizing supernova shock reflected from the old
boundary of the cavity. Alternatively, in accordance with the
model of Vidal et al. (2015), they might be a system of “il-
luminated” field lines on the back hemisphere of the loop
cavity.
–On several grounds we are not convinced by the suggestion
of Liu et al. (2014) that emission from the loop is seriously
contaminating derived CMB maps.
–We trace the magnetic field in Loop II for the first time. It fol-
lows the path of the loop, as in the other examples. We show
that the most prominent feature in the original definition of
Loop II is actually unrelated, and so the loop is somewhat
smaller and fainter than previously believed.
–The proposed extension to Loop III in the south suggested by
Vidal et al. (2015) is largely an artefact in the WMAP maps.
–The South Polar Spur, previously identified as the low-ledge
of Loop II, is actually the high-ledge of a smaller struc-
ture. Just like for Loop I, there is a filament of cold material
(traced by H iand thermal dust) running just outside it, with
parallel magnetic field. From the H ivelocity the distance is
at most a few hundred pc.
–Loops V and VI (Milogradov-Turin & Uroševi´
c 1997) and
S1 (Wolleben 2007) are not physical structures; our map
shows that major components of the proposed loops devi-
ate radically from the expected path. These components are
independent features, with much smaller angular size than
the suggested loops.
–We find two examples of smaller and more distant loops, one
of which arcs around the Cyg X star-forming region. Both
are associated with polarized dust features.
–We observe a clear outline of narrow highly polarized in-
tensity around the northern Fermi bubble above the Galactic
centre. The spectral index appears to be much flatter (β=
−2.54 ±0.16) than other filaments, but similar to that found
for the microwave haze. This is indicative of a more ener-
getic population of electrons along these lines of sight. We
note that the Southern extension of the Fermi bubble extends
further south than Loop I, which probably rules out an asso-
ciation of Loop I with the Fermi bubbles.
–We have discovered examples of anti-correlation of Hαand
polarized synchrotron intensity. The most striking is a high
latitude Hαfilament that appears as a local minimum in
polarization. Because of the low density of ionized gase
we can rule out Faraday depolarization. The observed anti-
correlation appears to be real in that there is no synchrotron
emission emanating from the region of the Hαfilament. The
negative velocity of Hαand the longitude imply that the fil-
ament could be lying at a distance of at least 2 kpc.
–We correlated the polarized maps with the Commander AME
map. We found that there was no significant correlation
except for the Pegasus plume region, where a low level
(9.0±1.9%) of polarization was observed, consistent with
contamination from surrounding synchrotron radiation.
Although we are beginning to understand the diffuse emission
at WMAP/Planck frequencies, there is still a long way to go. In
intensity, the component separation is still a major difficulty. In
particular, we have not been able to constrain the synchrotron
spectral dependence, due to the degeneracy and confusion with
free-free, AME, and CMB. New data in the frequency range
2–15 GHz are needed to more fully characterize all of the low-
frequency foreground components. Experiments such as S-PASS
at 2.3 GHz (Carretti 2011), C-BASS at 5 GHz (King et al. 2014),
and QUIJOTE at 11–19 GHz (Rubiño-Martín et al. 2012b) will
improve the situation considerably.
In polarization, the emerging picture is that synchrotron
emission is by far the dominant polarized component below
50 GHz, with very little contribution from free-free and AME.
The diffuse synchrotron emission has now been measured with
good S/N for a large fraction of sky. However, there are still re-
gions where the S/N is low. More importantly, the detailed spec-
tral dependence of synchrotron emission is still unknown.
Acknowledgements. This paper is dedicated to the memory of the late Professor
Rodney Deane Davies CBE FRS and Professor Richard John Davis OBE, both
of whom contributed greatly to the Planck project. The Planck Collaboration ac-
knowledges the support of: ESA; CNES and CNRS/INSU-IN2P3-INP (France);
ASI, CNR, and INAF (Italy); NASA and DoE (USA); STFC and UKSA
(UK); CSIC, MINECO, JA, and RES (Spain); Tekes, AoF, and CSC (Finland);
DLR and MPG (Germany); CSA (Canada); DTU Space (Denmark); SER/SSO
(Switzerland); RCN (Norway); SFI (Ireland); FCT/MCTES (Portugal); ERC
and PRACE (EU). A description of the Planck Collaboration and a list
of its members, indicating which technical or scientific activities they have
been involved in, can be found at http://www.cosmos.esa.int/web/
planck/planck-collaboration. This research was supported by an ERC
Starting (Consolidator) Grant (no. 307209) and STFC Consolidated Grant
(no. ST/L000768/1). We have made extensive use of the HEALPix package and
the IDL astronomy library. This research has made use of the SIMBAD database,
operated at CDS, Strasbourg, France.
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Planck Collaboration: Planck 2015 results. XXV.
Appendix A: Simple ILC method
The well-known ILC methods return an image of a component
with a known spectrum in the presence of other components
whose spectra are either not explicitly considered (Bennett et al.
2003) or only known in a subset of cases (Remazeilles et al.
2011a,b). In Sect. 3we have the more limited aim of eliminating
one or two of the major components with known spectra at low
frequencies (CMB and free-free emission), and returning a map
reasonably close to the sky at 28.4 GHz with these components
omitted.
We model the maps as sums of components with spatially-
invariant and known spectra, plus noise and other unmodelled
foreground emission:
˜
Tjk =Mki Ai j +Njk +Ojk (A.1)
where the indices are ifor component, jfor pixel and kfor band.
Ai j is the amplitude. The mixing matrix Mki =fi(νk)/UkCik,
where fi(νk) is the spectral form in Rayleigh-Jeans brightness
temperature at the band reference frequency νk,Ukis the unit
conversion from Rayleigh-Jeans to thermodynamic temperature,
and Cik is the colour correction (Planck Collaboration IX 2014;
Planck Collaboration II 2016)15. For the CMB, component 0, we
have Mk0=1.
Our aim is to find the weights for the frequency maps that
satisfy the conditions
X
k
Mkiwk=0 (A.2)
for the components to be eliminated, subject to the normalization
condition
X
k
Mk3wk=M03 (A.3)
where band k=0 is 30 GHz and component i=3 corresponds to
aβ=−3 power law, roughly appropriate for synchrotron emis-
sion, but also quite close to the effective spectrum of spinning
dust. We omit maps with frequencies in the range 50–120 GHz,
15 Colour corrections and unit conversions for WMAP were derived in
the same way as for the LFI, using the released WMAP bandpasses.
Colour corrections were based on the spectral index of the model eval-
uated at νk, except that we used the explicit dust colour corrections pro-
vided by the UC_CC code for the HFI bands (Planck Collaboration IX
2014).
where the foreground components are weak and their spectrum
(especially the dust components) is not well defined. The weak-
est point of our modelling is the assumption of a fixed spectral
shape for the thermal dust, but this component is very weak be-
low 50 GHz so that errors in the dust spectrum there have little
effect; the spectrum is mainly used in the interpolation from 353
to 143 GHz needed to generate a dust-cleaned CMB template.
We solve Eqs. (A.2) and (A.3) by restricting the analysis to
nfrequency maps, where ncomponents are dealt with, to make
Mki square, and then
w=M−1
0
.
.
.
M03
.(A.4)
In our final application, we have more bands than constraints and
so we found the weights by maximizing the S/N for our reference
component:
S/N=
X
k
wkMk3
,
X
k
w2
kσ2
k
1/2
,(A.5)
where, to keep the weights spatially uniform, we take σ2
kas the
median pixel variance for band k. We split the bands into a set k0
whose weights are varied by the fitting routine (AMOEBA in IDL),
and a set kwhose weights are found by linear regression. Our
constraint is now written
X
k
Mkiwk=δi3M03 −X
k0
Mk0iw0
k.(A.6)
Each AMOEBA iteration provides a trial set of w0
kvalues that are
used to find wkvia
wk=X
i
M−1
ki
δi3M03 −X
k0
Mk0iwk0
,(A.7)
and the full set of weights is then used to evaluate the S/N.
A25, page 45 of 45
